MAXI J1659−152: the Shortest Orbital Period Black-Hole Transient in Outburst

A&A 552, A32 (2013) Astronomy DOI: 10.1051/0004-6361/201219447 & c ESO 2013 Astrophysics

MAXI J1659−152: the shortest orbital period black-hole transient in outburst

E. Kuulkers1, C. Kouveliotou2,T.Belloni3, M. Cadolle Bel1,J.Chenevez4, M. Díaz Trigo5,J.Homan6, A. Ibarra1, J. A. Kennea7, T. Muñoz-Darias3,8, J.-U. Ness1,A.N.Parmar1,A.M.T.Pollock1, E. P. J. van den Heuvel9, and A. J. van der Horst9,10

1 European Space Astronomy Centre (ESA/ESAC), Science Operations Department, 28691 Villanueva de la Cañada, Madrid, Spain e-mail: [email protected] 2 Astrophysics Oﬃce, ZP12, NASA/Marshall Space Flight Center, Huntsville, AL 35812, USA 3 INAF-Osservatorio Astronomico di Brera, Via E. Bianchi 46, 23807 Merate (LC), Italy 4 National Space Institute, Technical University of Denmark, Elektrovej 327, 2800 Kongens Lyngby, Denmark 5 ESO, Karl-Schwarzschild-Strasse 2, 85748 Garching bei München, Germany 6 MIT Kavli Institute for Astrophysics and Space Research, 70 Vassar Street, Cambridge, MA 02139, USA 7 Department of Astronomy & Astrophysics, The Pennsylvania State University, 525 Davey Lab, University Park, PA 16802, USA 8 School of Physics and Astronomy, University of Southampton, SO17 1BJ, UK 9 Astronomical Institute “Anton Pannekoek”, University of Amsterdam, Science Park 904, 1098 XH Amsterdam, The Netherlands 10 Universities Space Research Association, NSSTC, 320 Sparkman Drive, Huntsville, AL 35805, USA Received 19 April 2012 / Accepted 21 December 2012

ABSTRACT

MAXI J1659−152 is a bright X-ray transient black-hole candidate binary system discovered in September 2010. We report here on MAXI, RXTE, Swift,andXMM-Newton observations during its 2010/2011 outburst. We ﬁnd that during the ﬁrst one and a half week of the outburst the X-ray light curves display drops in intensity at regular intervals, which we interpret as absorption dips. About three weeks into the outbursts, again drops in intensity are seen. These dips have, however, a spectral behaviour opposite to that of the absorption dips, and are related to fast spectral state changes (hence referred to as transition dips). The absorption dips recur with a period of 2.414 ± 0.005 h, which we interpret as the orbital period of the system. This implies that MAXI J1659−152 is the shortest period black-hole candidate binary known to date. The inclination of the accretion disk with respect to the line of sight is estimated to be 65–80◦. We propose the companion to the black-hole candidate to be close to an M5 dwarf star, with a mass and radius of about 0.15–0.25 M and 0.2–0.25 R, respectively. We derive that the companion had an initial mass of about 1.5 M, which evolved to its current mass in about 5–6 billion years. The system is rather compact (orbital separation of 1.33 R), and is located at a distance of 8.6 ± 3.7 kpc, with a height above the Galactic plane of 2.4 ± 1.0 kpc. The characteristics of short orbital period and high Galactic scale height are shared with two other transient black-hole candidate X-ray binaries, i.e., XTE J1118+480 and Swift J1735.5−0127. We suggest that all three are kicked out of the Galactic plane into the halo, rather than being formed in a globular cluster. Key words. accretion, accretion disks – binaries: close – stars: individual: MAXI J1659-152 – X-rays: binaries

1. Introduction (LMXBs), i.e., binaries containing a low-mass (typically 1 M) companion orbiting a neutron star or a black hole. Transient X-ray sources have been extensively studied since the BHXBs are primarily discovered when they enter outbursts advent of X-ray astronomy. Most are in binary systems compris- characterised by increased X-ray luminosities, by a factor of at ing at least one compact object, which is either a neutron star or least 100 in a few days, and a variety of spectral and tempo- a black hole. Determining the nature of the compact object is not ral variability states. The most contrasted states are the so-called trivial, as a number of their X-ray characteristics are common to “hard” and the “soft” states. The former, originally was also de- both types of objects. Black hole X-ray binaries (BHXBs) are fined as a “low” state, owing to the representation of the spec- uniquely identified by a mass determination of the compact ob- trum by a power law with spectral index 1.4–2.0 up to several ject in excess of 3 M. The first two confirmed BHXBs were hundreds of keV. Conversely, the latter “soft”, or “high”, state Cyg X-1 (Webster & Murdin 1972; Bolton 1972)andLMCX-3 is characterised by a tenfold increase of the ∼2–10 keV flux. (Cowley et al. 1983); these systems were, however, persistent Several other states have been defined in the literature, depend- X-ray sources. The very first transient source with a confirmed ing also on the evolution of the combined spectral and tempo- − black hole, A0620 00 (McClintock & Remillard 1986), was dis- ral variability properties (for extensive reviews, see Homan & ∼ covered in 1975, when it reached intensities of 50 Crab (Elvis Belloni 2005; Belloni 2010; and McClintock & Remillard 2006; et al. 1975). To date there are 20 known BHXBs, of which 17 are Remillard & McClintock 2006). known to be transient (e.g., Remillard & McClintock 2006). The On 2010 September 25 08:05 UTC, the Swift/Burst Alert BHXB transients belong to the class of low-mass X-ray binaries Telescope (BAT) triggered on a source located roughly 17◦ above the Galactic plane. The source was initially designated Table 1 is available in electronic form at http://www.aanda.org as gamma-ray burst (GRB) 100925A (Mangano et al. 2010), Article published by EDP Sciences A32, page 1 of 20 A&A 552, A32 (2013) and was monitored with the Swift/X-Ray Telescope (XRT). Our team was involved in various target of opportunity (ToO) Interestingly, the source flux did not decline during the next observations taken with RXTE, Swift,andXMM-Newton.Here several hours, as is usually the trend with GRBs. This unusual we present a detailed account of the X-ray and UV light curves behaviour and the source location near the Galactic bulge in- of MAXI J1659−152, taken over the course of the 2010 outburst dicated that it might not be a GRB but a new Galactic source with these satellites and with MAXI. A preliminary account of (Kahn 2010). Later that day, the Monitor of All-sky X-ray Image the results can be found in Kuulkers et al. (2010b,d, 2012a)and (MAXI) team reported the detection of a new hard X-ray tran- Belloni et al. (2010). The MAXI, RXTE and Swift data have sient, MAXI J1659−152, whose position was consistent with already been (partly) described elsewhere, however, the focus GRB 100925A and which had brightened since 2010 September was more on the combined spectral and timing behaviour of 25 02:30 UTC (Negoro et al. 2010). The discovery initiated MAXI J1659−152 (Kalamkar et al. 2011;Kenneaetal.2011; several multi-wavelength observations (ground and space-based; Muñoz-Darias et al. 2011b; Shaposhnikov et al. 2012; Yamaoka e.g., van der Horst et al. 2010; Kuulkers et al. 2010b,c; see also et al. 2012). We find that MAXI J1659−152 is indeed viewed at Kuulkers et al. 2012a). a high inclination. It still has the shortest orbital period known The Galactic origin of MAXI J1659−152 was also suggested to date, and is, therefore, a rather compact BHXB transient. We the next day from a combined UV-X-ray spectral-energy dis- present in Sect. 2 a detailed description of the data sets used in tribution analysis using data from the Swift/XRT and Swift/UV our analysis, and in Sect. 3 the results of our timing analysis of Optical Telescope (UVOT; Xu 2010). Finally, optical spec- the variations of the source light curve. We discuss the observed troscopy data taken with the ESO/Very Large Telescope (VLT) light curve dips in Sect. 4 and provide an interpretation on the X-shooter instrument showed various broad emission lines from characteristics of the binary system and its distance. the Balmer and Paschen series of H and of He II, as well as Ca II and Na I absorption lines from the interstellar medium, all at red- shift zero, clinching the Galactic nature of the source. Moreover, 2. Observations the double-peaked profiles of the emission lines indicated that In this paper we only concentrate on the light curves of the the source was an X-ray binary (de Ugarte Postigo et al. 2010; various instruments onboard MAXI (Matsuoka et al. 2009), Kaur et al. 2012). The binary nature of the source has by now RXTE (Bradt et al. 1993), Swift (Gehrels et al. 2004)and been established in several studies of its spectral and temporal XMM-Newton1 (Jansen et al. 2001). We refer to Table 1 for properties. Kennea et al. (2010) reported frequent intensity drops an observation log of the latter 3 satellites. The XMM-Newton in the X-ray light curve, attributed possibly to eclipses by a com- spectral analysis is deferred to a future paper. For a spectral panion star. Kuulkers et al. (2010d, 2012a) established a period analysis of the RXTE and Swift data we refer to Muñoz-Darias ∼ of 2.42 h using the X-ray Multi Mirror mission (XMM-Newton) et al. (2011b) and Kennea et al. (2011), respectively, as well as data, which was later confirmed by Belloni et al. (2010, Rossi Yamaoka et al. (2012). All our light curves have been subjected X-ray Timing Explorer (RXTE); also Kuulkers et al. 2012a) to a barycentric correction using the standard tools available. and Kennea et al. (2011; Swift). Kuroda et al. (2010)havere- ported optical variations up to 0.1 mag, consistent with a double- peaked modulation at a period of 2.4158 ± 0.0003 h. We con- 2.1. XMM-Newton cluded that this is the shortest BHXB orbital period measured We used Science Analysis System (SAS) version 11.0.0 together as yet (Kuulkers et al. 2010d, 2012a), and suggested that the with the latest calibration files to analyse the XMM-Newton data. source was most likely viewed at a high inclination (Kuulkers The European Photon Imaging Camera (EPIC)-MOS (Turner et al. 2012a). et al. 2001) cameras were not used during the observation in The source was identified as a BHXB through the fast tim- / order to allocate their telemetry to the EPIC-pn camera (Strüder ing behaviour observed with the RXTE Proportional Counter et al. 2001) and to avoid full scientific buffer in the latter. The Array (PCA), which was similar to that seen in stellar-mass EPIC-pn was used in timing mode, whilst the Reflection Grating black-hole transients (Kalamkar et al. 2011; Kennea et al. 2011; Spectrometer instruments (RGS1 and RGS2; den Herder et al. Muñoz-Darias et al. 2011b; Shaposhnikov et al. 2012; Yamaoka 2001) were operating in the standard spectroscopy mode. Due to et al. 2012). Estimates of the mass of the compact object range the brightness of the source the EPIC-pn data were affected by from 2.2–3.1 M (Kennea et al. 2011), to 3.6–8.0 M (Yamaoka ± pile-up (see below). Similarly, the count rate was also above the et al. 2012), to 20 3 M (Shaposhnikov et al. 2012). Part of limits of pile-up for the RGS2, for which the read-out is slower the discrepancy can be resolved by taking into account the by a factor of two compared to that of RGS1 since August 2007 spin of the black hole (Kennea et al. 2011; Yamaoka et al. − (see Sect. 3.4.4.8 of the XMM-Newton Users Handbook). 2012). The source distance estimates range between 1.6 4.2 kpc Standard data reduction procedures (SAS tasks epproc and (Miller-Jones et al. 2011) and 8.6 kpc (Yamaoka et al. 2011). − rgsproc) were used to obtain EPIC-pn and RGS calibrated After the main outburst, MAXI J1659 152 continued to event files. We used the task epfast on the event files to correct show low-level activity (e.g., Kennea et al. 2011). The flux ffi ff − for charge transfer ine ciency (CTI) e ects seen in the EPIC-pn of MAXI J1659 152 was observed to exhibit sudden intensity timing mode when high count rates are present. changes, e.g., during 2011 May, when Swift and Chandra de- The count rate in the EPIC-pn was close to, or above, the tected a reflare. The source then decayed again over a period of −1 ff ∼ 800 cts s level, at which X-ray loading and pile-up e ects be- 2 months (Yang & Wijnands 2011a,b,c; Jonker et al. 2012). It come significant. Pile-up occurs when more than one photon seemed to enter its quiescent state during the second half of 2011 is read in a pixel during a read-out cycle. This causes pho- (Yang & Wijnands 2011d; Russell et al. 2011; Kong et al. 2011). ton loss, pattern migration from lower to higher pattern types However, for a BHXB it was still too bright to be in true quies- and hardening of the spectrum, because the charges deposited cence (Jonker et al. 2012). A possible quiescent optical counterpart has been reported based on observations done on 2010 June 1 We note that the XMM-Newton observations were taken simultane- 19 (about 3 months before the start of the outburst) and 2012 ously with one of the INTEGRAL ToO observations of the source, see March 23 (Kong et al. 2010; Kong 2012). Kuulkers et al. (2012a).

A32, page 2 of 20 E. Kuulkers et al.: MAXI J1659−152 by more than one photon are added up before being read out Table 2. Log of observations with the XMM-Newton/OM. (see the XMM-Newton Users Handbook for more information on pile-up). Time1 Filter Flux Magnitude Since pile-up causes significant spectral distortion and a de- (ks) (10−15 erg cm−2 s−1 Å−1) cline in the count rate measured by XMM-Newton,weinves- 0.00 UVW1 1.82 ± 0.02 15.75 ± 0.02 tigated in detail its presence before extracting the time series. 4.52 UVW1 1.81 ± 0.02 15.76 ± 0.02 We used the SAS task epatplot, which utilises the relative ra- 9.05 UVW1 1.79 ± 0.02 15.77 ± 0.02 tios of single- and double-pixel events which deviate from stan- 13.57 UVW1 1.79 ± 0.02 15.77 ± 0.02 dard values in case of significant pile-up, as a diagnostic tool 18.09 UVW1 1.92 ± 0.03 15.69 ± 0.02 in the pn camera timing mode data and found that the spectrum 22.61 UVM2 1.18 ± 0.05 16.45 ± 0.05 was affected by pile-up. Next, we extracted several spectra se- 27.13 UVM2 1.32 ± 0.05 16.33 ± 0.04 lecting single and double timing mode events (patterns 0 to 4) 31.65 UVM2 1.22 ± 0.05 16.41 ± 0.05 but different spatial regions for the source. Source events were 37.97 UVM2 1.35 ± 0.05 16.31 ± 0.04 ± ± first extracted from a 62 (15 columns) wide box centred on 42.50 UVM2 0.95 0.03 16.69 0.03 the source position (Region 1). Next we excluded 2, 4, 6, 8 and Notes. (1) Start of the exposure in ks relative to UTC 2010 September 27 10 columns from the centre of Region 1 (Regions 2–6) and ex- 16:24:36 (=MJD 55 466.68375). tracted one spectrum for each of the defined regions. We found that the spectrum was free of pile-up after removing the central 10 columns. Then, we used this free of pile-up event list to ex- pointings, i.e., those with offset angles less than 0.02◦. The light tract the time series shown in this paper. curves were corrected for background, as estimated from the In the case of the RGS, we used Table 11 in the background model for bright sources. XMM-Newton Users Handbook to determine which CCDs were affected by pile-up. We found that CCDs 6, 7 and 8 from RGS2 2.3. Swift were above the pile-up limits and, therefore, we did not use them for analysis. We utilised the methods described by Evans et al. (2009)to The EPIC-pn and RGS time series at 1 s and 100 s time res- extract XRT (Hill et al. 2004; Burrows et al. 2005) X-ray olution were produced using the epiclccorr and rgslccorr light curves in the energy range 0.3–10 keV, with corrections ff tasks, respectively. for the e ects of pile-up, hot-columns and hot-pixels applied. In the EPIC-pn timing mode, there are no source-free back- These light curves, at 100 s time resolution, were extracted us- ground regions, since the point-spread function of the telescope ing the most accurate available localisation in the XRT coor- extends further than the central CCD boundaries. In the case of dinate system, derived from late-time PC mode data taken on RGS, since MAXI J1659−152 was very bright, its spectrum is 2011 February 6, i.e., 134 days after the initial detection of MAXI J1659−154, when the source was not affected by pile-up not significantly modified by the “real” background which con- / tributes less than 1% to the total count rate in most of the band- (Kennea et al. 2011). We used the XRT WT light curves between 2010 September 25 and October 22 (see Table 1), whenever the width. Therefore, we chose not to subtract the “background” ex- ff tracted from the outer regions of the central CCD (see also Done pointing o set was smaller than about 3.5 . This led to a total &DiazTrigo2010;Ngetal.2010). good-time exposure of about 123 ks. We used the BAT (Barthelmy et al. 2005) 15–50 keV light The final RGS light curve was calculated by combining curves (see Krimm et al. 2006)2 generated on 2011 October 25. RGS1 and RGS2, adding both orders 1 and 2. We did not use time bins with less than 500 s exposure time (see The optical monitor (OM, Talavera 2009; Mason et al. 2001) Kennea et al. 2011). was operated in EPIC imaging mode with the two filters UVW1 ∼ − ∼ − (λ 2500 3500 Å) and UVM2 (λ 2000 2600 Å). Ten con- 2.4. MAXI secutive exposures of 4200 s duration each were taken, first five in the UVW1, and then five more in the UVM2 filter. We derived The Gas Slit Camera (GSC; Mihara et al. 2011)ispartofthe the source magnitudes from the output of standard extraction us- payload of the MAXI mission, onboard the International Space ing the SAS task omichain. In Table 2, the average fluxes and Station (ISS). Only those data (version 0.3) when the instrument magnitudes in each filter are listed for each of the exposures. was operating at a high voltage of 1650 V were taken into account. In this paper we only focus on the light curves averaged 2.2. RXTE over an ISS orbit and on a daily basis, in the total 2–20 keV band3. RXTE monitored the source more than once daily throughout the outburst (see Kalamkar et al. 2011; Muñoz-Darias et al. 2011b; Shaposhnikov et al. 2012; Yamaoka et al. 2012). We used 3. Results 64 PCA (Jahoda et al. 2006) observations from 2010 September The main outburst light curve of MAXI J1659−152 has already 28 to November 8 (see Table 1); the total good-time exposure been shown in various papers, as well as a multitude of en- was about 138 ks. ergy range and time resolution combinations. Authors also used For the PCA data we used the ftools analysis suite (ver- different combinations of instruments. We refer to Kalamkar sion 6.9). We produced light curves from PCU2 with 16-s et al. (2011), Muñoz-Darias et al. (2011b) and Shaposhnikov bins over the full PCA energy range (using Standard 1 data), et al. (2012)fortheRXTE/PCA light curves of the averages i.e., 2–60 keV, and over three energy ranges, i.e., 2.1–4.9 keV, per observation, in various X-ray bands. Yamaoka et al. (2012) 4.9−9.8 keV and 9.8–19.8keV (using Standard 2 data). We applied the standard selection criteria for bright sources in our anal- 2 http://swift.gsfc.nasa.gov/docs/swift/results/ ysis. We included data when the elevation angle of the source transients/index.html above the Earth horizon was more than 10◦, and used only stable 3 http://maxi.riken.jp

A32, page 3 of 20 A&A 552, A32 (2013)

Fig. 1. Top panel: overview of the 2010/2011 outburst of MAXI J1659−152. a) MAXI/GSC and b) Swift/BAT daily light curves. Note that most of the measurements after day ∼100 are non-detections. In a) we mark the times of the Swift/XRT (“X”), Chandra (“Ch”) and radio (“R”) observations taken in 2011 (see Sect. 3.1). Middle panel: zoom-in from the top panel, focusing on the main part of the 2010/2011 outburst of MAXI J1659−152. c) RXTE/PCA PCU2 light curve, at a 16 s time resolution. The data within an observation are connected to guide the eye; errors are not plotted since they are negligible. The time-span of the XMM-Newton observations, discussed in this paper, is indicated in the top-left (“XMM”). The times of the transition dips (“td”; see text) are marked by arrows. The time spans for the four main epochs described in the text (“A”, “B”, “C”, “D”) are indicated at the bottom of this graph. d) MAXI/GSC and e) Swift/BAT light curves, showing averages per satellite orbit. The vertical dotted line marks the time of the first MAXI/GSC detection (Negoro et al. 2010). f) Hardness values as a function of time. Hardness is defined as the ratio of the 0.5 day averaged count rates in the Swift/BAT 15–50 keV band to the MAXI/GSC 2–20 keV band. Bottom panel: zoom-in from the middle panel, focusing on a possible 3-day variation discussed in Sect. 3.1. g) MAXI/GSC and h) Swift/BAT light curves showing the averages over a satellite orbit, between days 15 and 31. We also show the results of a sinusoidal fit to describe the possible ∼3 day variation, plus a constant, linear and quadratic term to account for the longer-term trend, between days 18 and 28. The horizontal dotted lines in a), b), d), e), g) and h) correspond to the zero level.

A32, page 4 of 20 E. Kuulkers et al.: MAXI J1659−152 / showed the RXTE PCA data at a time resolution of 16 s in var- 25−27/9/2010 ious X-ray bands. Kalamkar et al. (2011) and Yamaoka et al. (2012) also show the one-day averaged MAXI/GSC light curves. Additionally, the latter authors provided the daily-averaged Swift/BAT data. Kennea et al. (2011) showed the Swift/BAT data on a satellite orbit time scale, together with the Swift/XRT rates at a 100 s time resolution. Both Kalamkar et al. (2011)and Shaposhnikov et al. (2012) provided daily hardness averages

based on the RXTE/PCA data, whilst Kennea et al. (2011)pro- (0.3−10 keV) vided hardness curves during the outburst using the Swift/XRT −1 data.

We here present an overview of the outburst behaviour, com- XRT cts s bining the information at soft and hard energies, as well as at various time resolutions. We first describe the overall outburst light curve and then focus on two independent intensity variations seen, i.e., “absorption dips” and “transition dips”. We subsequently present a timing study of the data taken during 0 50 100 150 the time period when absorption dips were seen. We use days 0.5 1 1.5 2 2.5 since MJD 55464 to describe the epoch of time. This date is Time (MJD − 55464) / / close to the MAXI GSC and Swift BAT triggers of the outburst Fig. 2. Swift/XRT 0.3–10 keV light curve during the first few days of the (MJD 55464.10 and 55464.34, respectively, see Sect. 1). outburst of MAXI J1659−152, when Swift observed the source during every satellite orbit (days 0.3–2.6). Dipping activity is clearly apparent on a regular basis. 3.1. Overall outburst light curve In the top panel of Fig. 1 we show the overall outburst light variations on a similar scale are seen within three observations curve of MAXI J1659−152, using daily averages, at relatively of that period (days 23.7, 24.0 and 26.1, see Sect. 3.2.2). The soft energies (2–20 keV; Fig. 1a) and hard energies (15–50 keV; last part of the main outburst (days 30.5–44.1; outburst epoch D) Fig. 1b). After a fast rise of a couple of days, MAXI J1659−152’s shows a rather smooth, but not linear or power-law/exponential soft intensity fluctuates by 20–30% on a daily basis (see also like (see below), decay. We note that the observations during out- Kalamkar et al. 2011), on top of a general slow decline in in- burst epoch D have only a relatively short duration (see Table 1); tensity (∼0.1 cts cm−2 s−1 per 10 days), up to about day 30. It they exhibit no strong variability like that seen in the earlier out- then shows an exponential-like decay (with a decay constant burst epochs. of ∼7 days) up to about day 100. In the hard energy band In Figs. 1d and e the MAXI/GSC and Swift/BAT light (Fig. 1b), MAXI J1659−152 reaches a peak in intensity within curves of the outburst are presented, integrated over a satel- 3 days after the start. It subsequently decreases in a somewhat lite orbit. We note that the Swift/XRT (and Swift/UVOT, see irregular fashion (but smoother than with respect to the soft en- Kennea et al. 2011) coverage was up to day 27.3 (see Table 1 ergy band), until about day 65 (see also Kennea et al. 2011). and Kennea et al. 2011). The difference in rise time to maxi- After days 100 and 65, MAXI J1659−152 falls below the detec- mum between the soft and hard energy band, noted above, can tion limits of MAXI/GSC and Swift/BAT, respectively. In Fig. 1a be clearly seen. This is also borne out by the hardness curve we have also indicated the times of the post main-outburst X-ray showninFig.1f: just after discovery, the outburst is hard, and observations with Swift and Chandra, as well as radio observa- then softens during the following week (see also Kennea et al. tions, as reported in the literature (Kennea et al. 2011;Yangetal. 2011,usingSwift/XRT). Near the end of the main outburst, 2011a,b; Yang & Wijnands 2011a,b,c,d; Miller-Jones et al. 2011; i.e., around day 35, MAXI J1659−152’s radiation hardens again Jonker et al. 2012). (see also Kalamkar et al. 2011; Shaposhnikov et al. 2012,using In the middle panel of Fig. 1 we zoom in on the main part RXTE/PCA). The end of the main outburst is well covered by of the outburst, from just before the start of the outburst up to the MAXI/GSC (see also above), and, as noted earlier, it shows a the part when MAXI J1659−152 was too close to the Sun to be smooth, power-law/exponentiallike decay in the 2–20 keV band. observed during dedicated pointed observations with RXTE and The hardening around day 35 and the fact that RXTE/PCA is Swift. sensitive to hard (20 keV) X-rays causes the RXTE/PCA light Figure 1c shows the outburst as seen by the RXTE/PCA in curve to deviate from the power-law/exponential like decay de- its whole sensitive energy band (2–60 keV), at a 16 s time reso- scribed at softer X-rays (20 keV). lution. The variability during an observation in the first part of Between about days 18 and 28, the MAXI/GSC and the main outburst and up to before the soft X-ray peak of the Swift/BAT intensities seem to modulate on a several day time main outburst (i.e., up to about day 10, which we denote out- scale (Figs. 1d and e). Before and after this time period this is burst epoch A) is mainly due to periodic dips in the light curves, not evident. A closer look at the MAXI/GSC and Swift/BAT light which will be described in more detail in the next subsection, curves (Figs. 1g and h) shows a possible periodic variation on an Sect. 3.2.1. From before the peak of the outburst up to about about 3-day time scale, for about 3 cycles. Indeed, a pure lin- day 22 (outburst epoch B), variability within an observation is ear trend does not describe the data well in this time period: 2 = 2 = still seen, albeit less strong. This variability is due to flaring dur- χred 3.3 for 81 degrees of freedom (dof) and χred 2.1for ing the observations (see, e.g., Fig. 5c). During outburst epoch B 123 dof, for the MAXI/GSC and Swift/BAT light curves, respec- the average flux varies between 15–20% from observation to ob- tively. Adding a quadratic term does not significantly improve 2 = servation (see also Kalamkar et al. 2011). On days 22–31.5 (out- the situation (χred 3.4 for 80 dof and 1.8 for 122 dof, re- burst epoch C), apart from a general decreasing trend, the inten- spectively). A sinusoidal plus constant, linear and quadratic term sity appears to be varying between a low and a high value. Flux does significantly improve the fit, although it is still not ideal

A32, page 5 of 20 A&A 552, A32 (2013)

a 27−28/9/2010 −1

cts s RXTE/PCA (2−60 keV) 400 500 b −1 cts s

150 200 Swift/XRT (0.3−10 keV) c −1 cts s 60 80 100 RGS (0.3−2 keV) d −1 cts s

100 150 EPIC−pn (0.2−15 keV)

e EPIC−pn (2−10 keV/0.2−2 keV) 0.3 0.35 hardness Fig. 3. a) RXTE/PCA, b) Swift/XRT, c) XMM-Newton/RGS, d) XMM-Newton/EPIC-pn count rate, e) EPIC-pn hardness (ra- UVM2 tio of count rates in the 2–10 keV to 0.2–2 keV bands), and OM f) XMM-Newton/OM (UVW1 and UVM2) ﬂux curves on days

flux 2.7 to 3.3. The time resolution is 16 s for the RXTE/PCA data UVW1 and 100 s for the data from the other X-ray instruments. The 1 1.5 2 f X-ray data points are connected for clarity. The OM ﬂux is 2.8 3 3.2 in units of 10−15 erg cm−2 s−1 Å−1, see also Table 2. Note that UVW1 and UVM2 do not cover the same wavelength range Time (MJD − 55464) (see Sect. 2.1).

2 = 2 = (χred 2.2 for 77 dof and χred 1.3 for 119 dof, respectively). EPIC-pn light curves show shallower dip activity at half the re- The modulation may thus be not purely sinusoidal. Assuming currence time (see, e.g., Figs. 3c and d near day 3.15). The depth the signal is real, we derive a period of 3.15 ± 0.05 days (uncer- of the dips varies between 90% and 50% of the average out- tainty is determined by using Δχ2 = 1) in the Swift/XRT light of-dip-interval intensity in the light curves of Swift/XRT and curve and 3.04 ± 0.09 days in the Swift/BAT, i.e., consistent with XMM-Newton/RGS and EPIC-pn (Figs. 2 and 3), extracted with each other. The phase of the sinusoid is about 0.25 days earlier a time resolution of 100 s. for the Swift/BAT with respect to that of the MAXI/GSC. In Fig. 4 we show two epochs of the dip activity at a higher time resolution of 1 s as seen by the XMM-Newton/EPIC-pn. The 3.2. Recurrent intensity variations two epochs are one dip-cycle apart, i.e., about 0.1 day. Clearly, the dip morphology changes from cycle to cycle. Fast dipping 3.2.1. Absorption dips activity is observed, which can last up to 30 min. These fast dips Periodic drops of intensity, or dips, are observed shortly af- have a duration of 30 s, and often as small as 1 s. Occasionally, ter the start of the outburst, at day 0.3, up to day 8.2, first before and/or after the fast dips the intensity reaches the per- in the Swift/XRT light curves (e.g., Fig. 2, days 0.3–2.6, sistent level seen outside the dipping intervals. During the fast Fig. 3b, day 3.3), then in the XMM-Newton/RGS and EPIC- dips the intensity can drop down to about 15% of the persistent pn light curves (Figs. 3c and d, days 2.7–3.3), and finally in intensity, indicating that the shallower dips observed in Fig. 3 the RXTE/PCA curves (e.g., Fig. 3a, days 3.0–3.4 and Fig. 5a, are a consequence of averaging a smaller number of fast dips in day 7.1). A clear recurrence time of 0.1 days is best observed coarser time bins. in the XMM-Newton/RGS and EPIC-pn light curves, thanks to Next, we examine the hardness values of the the continuous coverage. In addition to the periodic dips, a lin- XMM-Newton/EPIC-pn (Fig. 3e) and RXTE/PCA (see Fig. 5) ear rise in the out-of-dip intensity is observed in the Swift/XRT light curves. The hardness ratio is defined as the ratio of the light curve (Fig. 2; see also Kennea et al. 2011), which ex- count rates in the 2–10 keV to the 0.6–2 keV bands for the tends throughout the XMM-Newton observations (Figs. 3cand EPIC-pn data, and 4.9–9.8 keV to the 2.1–4.9 keV bands for d) and is consistent with the MAXI/GSC light curve (Fig. 1d), the PCA data. During the dips the source hardens. This colour where the source reached its first plateau at soft X-rays, 20 keV, behaviour during dips is also confirmed by the Swift/XRT data around day 4 (i.e., just after the XMM-Newton observations). (see Kennea et al. 2011). We observe in particular a stronger The dips show irregular structure which lasts between hardening as the dip becomes deeper. As the out-of-dip intensity about 5 and 40 min. Occasionally, the XMM-Newton/RGS and increases, the dipping becomes shallower and the changes in

A32, page 6 of 20 E. Kuulkers et al.: MAXI J1659−152

structure time scale. Therefore, we do not discuss the UV data any further.

3.2.2. Transition dips On three occasions (days 23.7, 24.0 and 26.1; see Fig. 6), the X-ray light curves show further pronounced variations. The −1 RXTE/PCA light curve on day 23.7 (Fig. 6a) resembles that seen

cts s during observations with absorption dips. However, the hardness (Figs. 6d, g) correlates with the X-ray intensity, in contrast to the absorption dips. The RXTE/PCA light curves on days 24.0 and 26.1 (Figs. 6b and c) show the presence of two intensity levels which differ by about 30%, between which the source fluctu- a ates. The transitions are fast, i.e., they occur within tenths of 27/9/2010 seconds. The time spent in the upper level is about 1–2 min on 0 100 200 2.88 2.89 2.9 day 24.0 and at least 6–14 min on day 26.1. The time spent in the Time (MJD − 55464) lower level is about 12 to at least 20 min on day 24.0 and about 26 min on day 26.1. These light curves have been referred to as “flip-flop” light curves (see Kalamkar et al. 2011). However, the intensity fluctuations correspond to fast source state transitions (see Sect. 4.1.2), so we refer to them as “transition dips”. The intensity variations seen during the transition dips are more or less of the same order as the variations in the average intensity from observation to observation in outburst epoch C (see Fig. 1 and

−1 Sect. 3.1; see also Kalamkar et al. 2011). Contrary to Kalamkar et al. (2011), we ﬁnd clear hardness changes during the transi- cts s tion dips light curves (Figs. 6e and f). The hardness behaviour is similar to that seen on day 23.7, again opposite to that seen during the absorption dip episodes, i.e., the source is harder at the upper intensity level, and softens when the source transits to the lower level (Figs. 6handi). 5 b 27−28/9/2010 Although the phasing of the transition dips on day 26.1

0 100 200 is consistent with the phasing of the absorption dips (see 2.97 2.98 2.99 3 Sect. 3.2.1), the phasing of the transition dips on days 23.7 and Time (MJD − 55464) 24.0 is clearly not. Observations in between days 24.0 and 26.1 Fig. 4. Zoom in from Fig. 3dontheXMM-Newton/EPIC-pn do not show any dipping behaviour (even between phase 0.4 and (0.2−15 keV) light curves around the time of dip activity around day 0.6), and indicate that the transition dips do not recur with the 2.9 a) and 3.0 b), respectively. The time resolution is 1 s. All data with 0.1 day time scale. a fractional exposure of more than 0.5 per bin are shown. 3.3. Timing analysis of absorption dip activity the hardness ratio less pronounced (see Figs. 5a and b). This The Swift/XRT, XMM-Newton/EPIC-pn and RGS, and is shown in more detail in Fig. 5g, where we plot the changes RXTE/PCA data indicate that the absorption dips occur of hardness ratio as a function of count rate. The shape of the regularly, i.e., about every ∼0.1 day (Sect. 3.2.1), from the start curve traced by the PCA data from day 7.1 to 9.1 to 12.7 is of the main outburst up to about day 10 (outburst epoch A, see remarkably similar to the shape of the curve shown in Fig. 4 of Sect. 3.1). In LMXB dippers (Sect. 4.1.1) absorption dips recur Boirin et al. (2005) for the classical LMXB dipper 4U 1323−62 with the orbital period. Using our rather large time base line we as the source moves from deep dipping to shallow dipping and can establish an accurate period of the recurring dip activity in finally to a persistent state4. Using the absorption dip activity MAXI J1659−152. ephemeris from Sect. 3.3, we find that the dips recur at phase We performed a Lomb-Scargle (LS; Lomb 1976; Scargle 0.4–0.6 (see, e.g., Figs. 5a–f). 1982) period search, as well as a Phase Dispersion Minimisation Since the dip morphology and hardness behaviour resembles (PDM; Stellingwerf 1978) period search, on our data sets taken the dip phenomenon encountered in various high-inclination during outburst epoch A. The datasets of each instrument were X-ray binaries (Sect. 4.1.1), we refer to these dips as “absorp- treated separately. The LS and PDM searches were done over tion dips”. the period range 0.01–0.5days, with a frequency interval of We note that UV data simultaneous to the XMM-Newton/ 0.001 day−1. For the PDM search we used 20 phase bins with EPIC-pn data are available from the XMM-Newton/OM (see a phase bin width of 0.056. In the LS periodogram a peak in- Fig. 3f). Although the UV is variable, there does not seem to be a dicates a dominating period in the data set, whilst in a PDM correlation with X-ray intensity, in particular with the X-ray dip- periodogram a minimum indicates a dominating period. ping. Unfortunately, however, the time resolution is too coarse to investigate in detail the UV light curve on the X-ray dip 5 By folding the data on the recurrence time of the absorption dip activity, 0.1 day (see Sects. 3.2.1 and 3.3). 4 Note that the differences in the numbers are due to the differ- 6 The choice of number of phase bins and phase bin width is rather ent instruments used for the curves, RXTE/PCA in Fig. 5gand arbitrary; tests with various numbers of phase bins and different phase XMM-Newton/EPIC-pn in Fig. 4 of Boirin et al. (2005). bin width yielded consistent results.

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a 2/10/2010 b 4/10/2010 c −1 cts s 400 600 800 7/10/2010 e f hardness

0.6d 0.8 0.5 0.6 0.4 0.5 0.6 0.7 0.9 1 1.1 phase g hardness 0.6 0.8

300 400 500 600 700 800 cts s−1 Fig. 5. Intensity (top) and hardness ratio (middle) curves during three RXTE/PCA observations (day 7.1: a), d), respectively; day 9.1: b), e),respectively; day 12.7: c), f), respectively), using PCU2 data at a time resolution of 16 s. The light curves are for the 2–60 keV band, whilst the hardness is defined as the ratio of the count rates in the 4.9–9.8 keV and 2.1–4.9 keV bands. The curves are folded on the absorption dipping activity ephemeris (see Sect. 3.3). Bottom: hardness versus intensity diagram (HID) for the same three observations. Day 7.1: filled circles; day 9.1: filled triangles; day 12.7: filled squares).

The error on a period found was computed by construct- the instruments, and 0.001, 0.001, 0.002 and 0.007 for the PDM ing 1000 synthesised data sets. These data sets were obtained values of the RGS, EPIC-pn, XRT and PCA, respectively. We by distributing each data point around its observed value, by an refer to this level as the noise level. amount given by its error bar multiplied by a number output by Another way to characterise the uncertainty in the period, a Gaussian random-number generator with zero mean and unit which is more conservative, is to use the width of the peak or variance. The measured standard deviation of the positions of minimum of the periodograms, e.g., the half-width at half maxi- the deepest minima in Θ or highest peak in power in the re- mum or minimum (HWHM). sulting periodograms was taken as the error. These latter peri- The results for the different instruments are discussed in the odograms were done in a narrow range around 0.1 day, i.e., next subsections, and a summary of the best-found periods near between 0.0909 and 0.1111 days, with a frequency interval of 0.1 days with their associated errors is given in Table 3. We find, 0.00009 days. that the HWHM values are a factor of 25–200 times larger than To check the significance of the peaks and minima found in those derived by the measured standard deviations of the po- the LS and PDM diagrams, respectively, we randomised the data sitions of the peaks or minima, as described above. Since the and calculated the resulting periodogram. The randomisation spread in the best-found periods is of the order of the HWHM was done by keeping the time tags and randomly distributing the values, we use these values as a final indicator of the uncertainty intensities of the data sets which were input to the periodogram in the derived periods. programs. This was repeated 1000 times and the resulting averaged periodogram was used to evaluate the significance of the 3.3.1. XMM-Newton/EPIC-pn and RGS peaks/minima. We find that the values of the power (LS) or amplitude Θ (PDM) are narrowly distributed around 1 for all pe- We used the EPIC-pn and the RGS data with 100 s time resolu- riods investigated: the variances are 1 for the LS values of all tion. The highest peak in the LS periodograms is at a period near

A32, page 8 of 20 E. Kuulkers et al.: MAXI J1659−152

b 19/10/2010 c 21/10/2010 −1 cts s 300 400 500 a 18/10/2010 d e f hardness 0.4 0.5 0.6 0.3 0.4 0.5 0.6 0.5 0.6 0.7 0.8 0.3 0.4 0.5 0.6 phase g h i hardness 0.4 0.5 0.6 300 400 500 300 400 500 300 400 500 cts s−1 Fig. 6. Intensity (top) and hardness ratio (middle) curves during the RXTE/PCA observations with transition dips (day 23.7: a), d), respectively; day 24.0: b), e), respectively; day 26.1: c), f), respectively), using PCU2 data at a time resolution of 16 s. The light curve is for the 2–60 keV band, whilst the hardness is deﬁned as the ratio of the count rates in the 4.9–9.8 keV and 2.1–4.9 keV bands. The curves are folded on the absorption dipping activity ephemeris (see Sect. 3). Bottom: hardness versus intensity diagram (HID) for the same three observations as above (g), h),andi) for day 23.7, 24.0 and 26.1, respectively), using the same data.

Table 3. Results of the PDM and LS period searches in the light curves for the diﬀerent instruments.

PDM LS Instrument Period Error HWHM Period Error HWHM (energy) (days) (days) (days) (days) (days) (days) RGS (0.3–2 keV) 0.10070 0.00013 0.0042 0.100985 0.000098 0.0071 EPIC-pn (0.2–15 keV) 0.09931 0.00006 0.0049 0.100394 0.000190 0.0090 XRT (0.3–10 keV) 0.10054 0.00004 0.0010 0.100282 0.000011 0.0010 PCA (2–60 keV) 0.100582 0.000001 0.00022 – – – PCA (2–5 keV) 0.100590 0.000018 0.00021 – – – PCA (5–10 keV) 0.100590 0.000019 0.00023 – – – PCA (10–20 keV) 0.100580 0.000024 0.00022 – – –

0.1 day (top two left panels of Fig. 7), i.e., 0.1004 ± 0.0090 days respectively) clustered within 0.2 in phase space. For the other and 0.1010 ± 0.0071 days, respectively, for the EPIC-pn and two periods, the same absorption dip activity is distributed along RGS data sets. Three deep minima are visible in the RGS and all phases. This strengthens our conclusion that the fundamental EPIC-pn PDM periodograms, near 0.1, 0.2 and 0.4 days (top period is near 0.1 days, since absorption dip activity is expected two right panels of Fig. 7). The deepest minimum is not at the to occur at restricted orbital phases (see Sect. 4.1.1). best period found with the LS search, however. Inspection of the folded light curves on the three periods 3.3.2. Swift/XRT found with the PDM, reveals that only when folding the data on the 0.1 day period is the absorption dip activity (i.e., at A LS and PDM search on the XRT data during outburst epoch count rates 140 c s−1 and 75 c s−1 for the EPIC-pn and RGS, A did not reveal any signiﬁcant period, except for the satellite

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RGS Θ

power RGS 0.6 0.8 1 0 204060 EPIC−pn Θ

power EPIC−pn 0.6 0.8 1 0 204060

XRT Θ

power XRT 0.6 0.8 1 050

PCA Θ power

0.8 0.9 1 PCA 0 65 130 0.010.02 0.05 0.1 0.2 0.5 0.010.02 0.05 0.1 0.2 0.5 Period (days) Period (days)

Fig. 7. Results of the Lomb-Scargle (left) and Phase Dispersion Minimisation (right) period analyses (see text) of the XMM-Newton/EPIC-pn (top), XMM-Newton/RGS (second panel), Swift/XRT (third panel)andRXTE/PCA (bottom) datasets. orbital period around the Earth (see below). This we attribute To inspect whether the fundamental period is related to the to the variation in intensity of the overall main outburst light data sampling (such as due to the satellite orbit), we constructed curve, which is of the same order as the drops in intensity during a Fourier transform (FT) of the window function. This window the absorption dips (see Fig. 2 and Kennea et al. 2011). The in- function was determined by setting the intensities of the time crease in the out-of-dip intensity is rather gradual during the first series to zero. For the FT we used the same frequency settings 6 days. Thereafter, the out-of-dip light curve varies irregularly on as the LS and PDM searches. The largest peak in the FT pe- a time scale of days. We, therefore, first detrended the data using riodogram is evidently at the satellite orbit period around the a multi-order polynomial (see also Kennea et al. 2011); a third- Earth of about 0.067 days (Fig. 8, top). Thus our best period is order polynomial describes the overall out-of-dip light curve up not related to any peak in the FT periodogram, and can thus be to about outburst day 6 sufficiently well. Including data after considered to be intrinsic to the source (see also Kennea et al. day 6 contaminated our period search significantly, and had the 2011). effect of diminishing the peak and minima in our LS and PDM searches, respectively, near 0.1 day. We, therefore, use these data only up to day 6, in the remainder of this subsection. 3.3.3. RXTE/PCA The highest peak in the LS periodogram is at The overall variability of the PCA data (see Sect 3.1 and Fig. 1c), 7 0.1003 ± 0.0010 days (see Fig. 7, left panel) . Several minima prevented the LS and PDM periodograms to show any significant are found in the PDM periodogram (Fig. 7, right panel), peaks or minima, respectively, except at the satellite orbital pe- including those seen at the same periods as in the PDM peri- riod around the Earth (see below). Detrending the PCA data as odograms of the XMM-Newton data (Sect. 3.3.1). The minimum done for the XRT did not improve our period searches, however. at 0.1005 ± 0.0010 days is clearly not the dominant period in This is due to the fact that the PCA data sampling (as well as the this PDM search. However, inspection of the folded light curves energy range) is different from that of the XRT data. Moreover, on the various periods with peaks or minima in the LS and PDM there are significant day-to-day variations in the average out-of- searches, respectively, show that, again, only for the period near dip PCA intensity, which cannot be described by a simple poly- 0.1 day, the dip activity is clustered within 0.2 in phase space, nomial. We, therefore, renormalised the PCA light curves during as expected for absorption dips (see Sect. 4.1.1). the whole outburst epoch A in our search for periodicities. For We note that performing the period search after renormalis- each observation interval (generally corresponding to an RXTE ing the XRT data in a similar manner as we did for the RXTE satellite orbit) we determined the mean count rate. This value data (see Sect. 3.3.3) did not reveal significant power at the above was subtracted from the light curves corresponding to each of reported periods. This is because the XRT observations were in these observation intervals. / general shorter than the RXTE PCA, resulting in averages which The lowest minimum in the PDM periodogram is at a pe- are higher when there is a dipping period, and therefore the vari- riod of 0.10058± 0.00022 days (Fig. 7, bottom right). A PDM ations due to dipping are diminished. search in three energy bands (2–5 keV, 5–10 keV and 10–20 keV, see Sect. 2.2) shows the deepest minima at the same period (see 7 We note that the period reported by Kennea et al. (2011), 0.1008 ± 0.0037 days, was determined using the LS search on the de- Table 3). Again, the period quoted above does not coincide with trended first 12.8 days of the XRT/WT data. Their error in the pe- the strong peak at about 0.065 days in the FT of the PCA window riod is derived by fitting a Gaussian to the Lomb-Scargle periodogram function (Fig. 8, bottom panel). around the peak, and taking a value of 2.7σ,whereσ is the width of the A LS search on the PCA data did not reveal a peak near Gaussian. 0.1 day period, nor near any of the other periods found in the

A32, page 10 of 20 E. Kuulkers et al.: MAXI J1659−152

XRT window power −100 0 100 renormalized intensity

0 0.2 0.4 0 0.2 0.4 0.6 0.8 1 phase

Fig. 9. Detrended RXTE/PCA (2–60 keV) 16 s data folded on the PCA window ephemeris given in Sect. 3.3.4.

The folded detrended PCA curve during the dip epoch using power this ephemeris is shown in Fig. 9. The dipping activity has a duty cycle of about 0.2 in phase, around phase 0.5. The mid-point between start and end of the duty cycle corresponds to T0,dip = MJD 55 467.0904 ± 0.0005. 0 0.1 0.2 0.3 0.010.02 0.05 0.1 0.2 0.5 Period (days) 4. Discussion Fig. 8. Fourier transforms of the window functions using the 4.1. Origin of intensity variations RXTE/PCA (top)andSwift/XRT (bottom) data. The peaks in power density spectra are consistent with the satellite orbital periods around 4.1.1. Absorption dips the Earth, i.e., 0.065 days (94 min) and 0.067 days (96 min), respectively for the RXTE/PCA and Swift/XRT. We find regular absorption dips in the X-ray light curves during the outburst of MAXI J1659−152 between outburst days 0.3 and 8.2. Absorption dips recur at the orbital period of the system and are thought to be caused by obscuration by material located in a XMM-Newton and XRT period searches (Fig. 7, bottom left)8. thickened outer region (“bulge”) of the accretion disk due to its We attribute this to the highly non-sinusoidal nature of the mod- interaction with the inflowing gas stream from the companion ulation in the renormalised light curves at 16 s time resolution, (e.g., White & Swank 1982; Walter et al. 1982; see discussion and the possible imperfection of our method of renormalisation below). The presence of absorption dips allows a precise mea- for the LS search. surement of the orbital period and is a signature of high inclination (e.g., White & Swank 1982; White & Mason 1985; see also Díaz Trigo et al. 2009). 3.3.4. Recurrence period of dipping activity We determine the recurrence period of the dips to be The XMM-Newton and XRT LS period searches and the 2.414 ± 0.005 h (see also Kuulkers et al. 2012a;Kenneaetal. PCA PDM period search reveal one common best period, 2011). By analogy with other classical dippers, we identify i.e., ∼0.1 day, with peaks/minima in the periodograms well this period with the orbital period of a system. The fastest re- above/below the noise level. Inspection of the individual light volving binary was Swift J1753.5−0127 (3.2443 h; Zurita et al. curves of the various instruments, and the folding of the data on 2008;Durantetal.2009). As suggested by Kuulkers et al. the 0.1 day period leading to restricted phase range of dipping (2010d, 2012a), if the compact object in MAXI J1659−152 is (see below), supports the main period to be at that value. We at- indeed a black hole (Kalamkar et al. 2011; Kennea et al. 2011; tribute the peaks/minima at half this period in the XMM-Newton Muñoz-Darias et al. 2011b; Shaposhnikov et al. 2012; Yamaoka periodograms to the intermediate dipping episodes seen. The et al. 2012), its 2.4 h period is the shortest among the currently peaks/minima at multiple times the above quoted period are due known BHXB sample (see, e.g., Ritter & Kolb 2003)9. to the fact that the morphology of the dipping changes from cy- 9 − cle to cycle, as well as the fact that not all dipping periods are A possible exception may be Swift J1357.2 0933, based on an indi- rect estimate of the orbital period of 2 h by Casares et al. (2011). We sampled well enough. note, however, that the outburst amplitude in the optical (∼6 mag, Rau The RXTE/PCA provides the strongest constraint on the pe- et al. 2011) is incompatible with such a short period, based on the empir- riod. We, therefore, conclude that the fundamental period of ical relation between the outburst amplitude and orbital period (Shabaz dipping activity is at 0.10058 ± 0.00022 days. We arbitrarily & Kuulkers 1998), see also Sect. 4.6. Another short-period binary can- set the zero point of the absorption dip activity ephemeris didate, MAXI J1305−704, was put forward recently by Kennea et al. to the first data point of the first RXTE/PCA observation at (2012b). This new transient (Sato et al. 2012) showed absorption dips. ∼ day 3.0. Therefore, the absorption dip activity ephemeris is A periodicity of 1.5 h was reported (Kennea et al. 2012b), but this was T = MJD 55467.039561 + 0.10058(22) × E, where E is the put into doubt by Kuulkers et al. (2012b). It has been suggested that 0 MAXI J1305−704 is a BHXB based on outburst optical amplitude, blue cycle number of the period. optical spectral energy distribution and hard X-ray spectrum (Greiner et al. 2012; see also Kennea et al. 2012a), as well as the occurrence of 8 The highest peak in the LS periodogram is at 0.0335 days. However, a state transition (Suwa et al. 2012). However, these are features which folding the data at this period reveals the absorption dips to be dis- are seen in neutron star LMXBs as well (see Suwa et al. 2012; Kennea tributed along all phases. et al. 2012b).

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We constrain the inclination of MAXI J1659−152 to be be- other BHXBs undergoing an outburst, like 4U 1630−47 and tween about 65◦ and 80◦ from the presence of the periodic ab- GRO J1655−40 (Kuulkers et al. 1998, 2000; Tomsick et al. sorption dips, due to material in the line of sight that obscures up 1998). It is plausible that for transient BHXBs, changes in the to about 90% of the total emission at given cycles, and the ab- accretion mode cause the appearance or disappearance of dips. sence of eclipses. We base the lower limit on the size of the bulge Kuulkers et al. (2000) interpreted the (deep) absorption dips dur- (and not on the disk opening angle which has been generally esti- ing the rise and plateau phase of the outburst in GRO J1655−40 mated to be 12◦, e.g., de Jong et al. 1996; Bayless et al. 2010). as due to an absorbing medium which is filamentary in nature. White & Holt (1982) estimated the size of the bulge responsi- These filaments could be due to the stream of material coming ble for absorption dips as 19◦ ± 6◦ for the LMXB 4U 1822−37. from the companion star splashing into the accretion disk and Taking this as typical for LMXBs, we can thus set a lower limit overflowing above and below the impact region (e.g., Livio et al. on the inclination of 65◦. We note that, if the accretion disk is 1986; see Kuulkers et al. 2000, and references therein). If the tilted or warped, the lower limit for the inclination could be as inclination is high enough, the impact region itself comes also low as 55◦, taking into account that for generic LMXBs, a disk into the line of sight (e.g., Frank et al. 1987). However, the pres- tilt of about 10◦ is expected (Foulkes et al. 2010). An upper ence of absorption features all around the orbit for neutron stars limit for the inclination of 80◦ is derived from the absence of (e.g., Parmar et al. 2002) shows that at least part of the photo- eclipses (e.g., Horne 1985), the spectral type of the companion ionised plasma is distributed equatorially along the whole plane star (M5V, see Sect. 4.2.1), and the fact that the companion is of the disk, indicating that absorption is due to a structure in the filling its Roche lobe (see, e.g., Motch et al. 1987). Here we are disk rather than by filaments. In that scenario, the cause for the assuming that the source of emission is point-like (for an ex- disappearance of dips in BHXBs could be, e.g., a strong ioni- tended source one would not be able to see full eclipses for any sation of the plasma in bright (but hard) states of the outburst, of the LMXBs). which renders the plasma transparent and therefore invisible to The absorption dips in MAXI J1659−152 share many of the us. Alternatively, a change of the structure of the disk could di- properties of those seen in classical absorption dipping systems. minish the thickness of the bulge and cause the absorption dips They change from period to period, they are fast, and the obscu- to disappear. ration can be large, i.e., down to about 90% of the total intensity (see, e.g., White & Mason 1985; Parmar & White 1988). Boirin 4.1.2. Transition dips et al. (2005) and Díaz Trigo et al. (2006) were able to model the changes in both the narrow X-ray absorption features and the As noted by Kalamkar et al. (2011), during the second epoch continuum during the dips from all the bright dipping LMXBs of the outburst rapid and sharp flux variations – transition observed by XMM-Newton by an increase in the column density dips – were seen, resembling the “flip-flop” and “dip” light and a decrease in the amount of ionisation of a photo-ionised curves in, e.g., GX 339−4 (Miyamoto et al. 1991), GS 1124−68 absorbing plasma. The changes in the hardness ratio observed (Takizawa et al. 1997) and XTE J1859+226 (Casella et al. 2004). in the dips in MAXI J1659−152 are consistent with absorption In all these cases 10–20% changes in intensity were seen. by neutral and photo-ionised plasma, in the sense that they are We find that MAXI J1659−152 softens when the intensity de- energy dependent. A further support to the existence of neutral creases, in contrast to the hardening seen during the absorp- and photo-ionised plasma is the presence of various stages of tion dips. In the case of GX 339−4 (Miyamoto et al. 1991) dipping: persistent, shallow and deep dipping states. The fact and XTE J1859+226 (Casella et al. 2004) significant hardness that dips become shallower and less energy dependent as the changes could be discerned as well, with behaviour similar to count rate increases could be a consequence of the photo-ionised that seen for MAXI J1659−152. The clear difference in hardness plasma becoming more and more ionised and transparent as it is behaviour of the absorption and transition dips, as well as the illuminated by the X-rays of the central region. However, we fact that they occurred in well separated phases of the outburst, note that a definite confirmation of an increase of neutral and suggests that the two phenomena have a different origin. ionised plasma during dips for MAXI J1659−152 is only possi- The transition dips in MAXI J1659−152 occurred during the ble after spectral analysis. first soft excursion of the source (see Fig. 10). During this pe- Absorption dipping in other LMXBs occurs mainly around riod the count rate differences between consecutive observations orbital phase 0.7–0.9, where eclipses are expected at phase zero were of the same order as the count rate changes seen during the if we view the accretion disk edge-on, i.e., when the companion transition dips (with the softer observations having lower count star is closest to us and in front of the neutron star (e.g., Parmar rates), suggesting that additional transitions took place between & White 1988). Occasionally, absorption dips with a 0.5 phase observations. difference with respect to the phase at which “regular dips” oc- In the other sources in which transitions dips have been seen, cur are observed. These “anomalous dips”, or “secondary ab- the transitions were often accompanied by pronounced changes sorption dips”, were also seen in other dipping systems (such as in the power density spectra (Miyamoto et al. 1991;Takizawa XB 1916−053; e.g., White & Swank 1982; Walter et al. 1982; et al. 1997; Homan et al. 2001; Casella et al. 2004). The power Smale et al. 1988;Boirinetal.2004). They are explained as density spectra from the two observations that showed the transi- being due to material migrating to the other side after impact tions dips in MAXI J1659−152 (days 24.0 and 26.1) were not of with the disk, or the accretion stream partly freely overflow- high enough quality to detect significant changes in the power- ing the disk or bouncing of the disk rim and then overflow (see, density spectral properties. However, by analysing the averaged e.g., Frank et al. 1987; Armitage & Livio 1998, and references power density spectra from the combined high and low count therein). In the latter case the flow may impact the disk near rate levels during the first “soft excursion”, Kalamkar et al. the circularisation radius, either causing a second bulge (see, (2011) were able to see indications for an additional broad bump e.g., Frank et al. 1987; Armitage & Livio 1998, and references around 7–8 Hz in the power density spectra of the low count rate therein) or bouncing of the disk again (Kunze et al. 2001). selection. We note that the transitions in other sources often in- The absorption dips appear for only part of the outburst volve so-called “type B” quasi-periodic oscillations, QPOs, ei- in MAXI J1659−152. This has been observed already for ther in the low or high count rate level power density spectra.

A32, page 12 of 20 E. Kuulkers et al.: MAXI J1659−152

Fig. 10. Hardness-intensity diagram obtained using all the RXTE/PCA observations available (adapted from Muñoz-Darias et al. 2011b). Intensity corresponds to the count rate in the 2–15 keV band and hardness is deﬁned as the ratio of counts in 6.1–10.2 keV and 3.3−6.1 keV bands. Each point corresponds to the average over an observation. Observations with a star correspond to those with a so-called type-B QPO in the power-density spectra (see text). A solid line joins consecutive observations starting from observation with a big, open circle (top right). Observations taken after the last type-B QPO are joined by a dotted line. The three arrows mark the observations with clear transition dips on days 23.7, 24.0 and 26.1. All observations between days 0 and 10, i.e., when absorption dips occur, have hardness values 0.45.

The broad excess seen by Kalamkar et al. (2011) is too broad to and assumed that all values and parameters are independent. be identified as a type B QPO, and is more likely to be a peaked The spread in the resulting values was used as the uncertainty. noise component. We note that parts of the discussion in the following subsections Transition dips are most likely the result from instabilities in already appeared in Kuulkers et al. (2012a) and Kennea et al. the inner accretion flow (see Miyamoto et al. 1991,foranex- (2011). We here update various of the values, based on our more ample interpretation), but their exact origin remains unknown. refined analysis. We indicate where Kennea et al. (2011) derived The observations of MAXI J1659−152 do not provide signifi- comparable properties. cant new insights into the nature of these instabilities, but they do show that transition dips can also be found in states that are 4.2.1. Mass and radius of the companion star slightly softer than those in which they have been observed in other sources (i.e., states in which type B QPOs are observed). 4.2.1.1 Unevolved main sequence companions This trend is very clearly seen in Fig. 10, where the type B QPOs We use here only the observed empirical mass-radius relation- all occur within a hardness range of 0.36–0.40, whilst all transi- ships for donor stars in Cataclysmic Variable (CV) binaries: tion dips occur at hardness 0.36. (i) Using the linear relationship for CV secondaries by Smith Around the time of the occurrence of transition dips we find & Dillon (1998; their Eqs. (9) and (12)) we derive the mass of some marginal evidence for the soft and hard X-ray light curves the companion, M2 = 0.19 M, and its radius, R2 = 0.24 R, to modulate on a 3 day period. We speculate that this period with an estimated uncertainty of 0.05 M and 0.02 R, respec- may be related to a disk precession period. Systems with ex- tively. (ii) Using the fact that the orbital period of 2.4 h of treme mass ratio’s (i.e., q 0.33), like MAXI J1659−152, are MAXI J1659−152 puts it right in the so-called “period gap” of vulnerable to a 3:1 orbital resonance within the accretion disk. CV-binaries, the empirical mass-radius relation for CV secon- This causes the disk to be eccentric and to slowly precess on daries of Knigge et al. (2011; their Fig. 4) yields a very precise time scales of days to weeks, which may be discernable in the mass estimate of M2 = 0.20 M,andaradiusofR2 = 0.26 R light curves (Whitehurst 1988; Whitehurst & King 1991; Lubow with a rather large uncertainty of 0.04 R. Since both estimates 1991a,b; see also Haswell et al. 2001). Periodic variations are are fully consistent with each other, and since the mass estimate also foreseen in this model with a period slightly longer than the from Knigge et al.’s relation is very precise, we will adopt the orbital period. However, in LMXBs this phenomenon is inclina- latter values as the real parameters of the companion star in the tion dependent: in systems with a high orbital inclination only system. The mass and radius correspond to spectral type M5V orbital modulations due to the heated face of the companion star if the star would be on the main sequence. However, it is well are expected (Haswell et al. 2001). This is consistent with the known that secondaries in CVs and LMXBs are often not pre- fact that the X-ray and optical light curves show the same (or- cisely on the main sequence, and tend to be a bit overluminous bital) period (see Sect. 1). The transition dips occurred during the for their masses. The fact that the companion here is still filling time of the 3 day modulation. Possibly, the non-axisymmetric its Roche lobe (as it is transferring mass), whilst its orbital period accretion disk modulates the inner accretion flow, giving rise to is located in the period-gap of CVs implies that the companion the sporadic transition dips in the X-ray light curves. is not a normal main sequence star, since in that case it would at that period not be filling its Roche lobe. It must therefore be a nuclearly somewhat evolved star (see below). 4.2. Binary system and evolutionary state

Using various existing empirical relationships one can get an 4.2.1.2 Nuclearly evolved stripped companions idea of the dimensions of the system. To derive uncertainties on the resulting values, where appropriate we randomly distributed If the companion star was originally much more massive and the observed values and equation parameters around their values became a nuclearly evolved star, it will now be a stripped using a Gaussian distribution with width equal to their errors, evolved star. Such a star will be He-rich and will obey a diﬀerent

A32, page 13 of 20 A&A 552, A32 (2013) mass-radius relation. We argue here that the companion is in- have orbital angular momentum loss by magnetic braking, the deed such a star and started out its life with a mass between companion star must have a convective envelope (Verbunt & 1.0 and 1.5 M, and underwent considerable nuclear evolution. Zwaan 1981; Rappaport et al. 1983; Knigge et al. 2011). This The reasons for this are as follows. Several authors, starting with implies for main-sequence stars, that the initial mass cannot have Pylyser & Savonije (1988), made calculations of the evolution of been larger than 1.5 M, since above this mass main-sequence LMXBs driven by the internal evolution of the companion star stars have radiative envelopes. The conclusion is, therefore, that in combination with orbital angular momentum loss by “mag- the companion star in MAXI J1659−152 must have started out netic braking” and gravitational waves (see Pfahl et al. 2003, with a mass between 1.0 and 1.5 M. For examples of the evolu- for recent evolutionary calculations for LMXBs). This evolution tion of systems with decreasing orbital periods, and with initial leads to such stripped evolved companion stars. The reason why companion masses of 1.0 or 1.5 M and a 4.0 M black hole ac- the companion of MAXI J1659−152 must be such an evolved cretor, see the evolutionary sequences 50 and 55 (Table 1) of star is that its orbital period of ∼2.4 h falls right in the middle Pylyser & Savonije (1988). The sequences Z55 and A55 in this of the “period gap” of CVs. In CVs with a normal H-rich main- paper started out with a companion mass of 1.5 M and evolved sequence companion the mass transfer is driven solely by orbital to an orbital period of 2.4 h, where the companion masses had angular momentum loss, due to magnetic braking and gravita- become 0.19 and 0.17 M, respectively, with a central H con- tion radiation (see, e.g., Howell et al. 2001; Knigge et al. 2011). tent of 0.21 and 0.09, respectively. The ages of the systems are Such systems evolve to shorter and shorter orbital periods; when then about 4.6 and 5.7 billion years, respectively, whilst the mass −11 −1 they arrive at the upper edge of the period gap, at Porb 3h,they transfer continues at a very low rate (below 10 M yr ). It stop transferring matter, because magnetic braking stops (Spruit appears that these sequences may well represent the evolution of & Ritter 1983;Howelletal.2001; Knigge et al. 2011). The com- MAXI J1659−152. panion star which was somewhat out of thermal equilibrium then shrinks, the orbit also shrinks – though slower – due to gravita- 4.2.2. Size of the X-ray source and of the absorbing object tional wave losses. The mass transfer resumes only when the system has reached the lower edge of the period gap at about Both Kennea et al. (2011) and Yamaoka et al. (2012)derive Porb 2 h. This type of evolution will hold for all binaries with a mass of the black hole of MBH 2.2–3.1 M and MBH a low-mass main-sequence companion, regardless whether the 3.6−8.0 M, respectively, based on the minimum innermost disk accretor is a white dwarf or a black hole. The fact that the or- radius from X-ray spectral fits being the innermost stable circu- bital period of MAXI J1659−152 falls in the middle of the pe- lar orbit (assuming a distance of >6.1 kpc and 5.3–8.6 kpc, re- riod gap, therefore, indicates that its companion is not a nor- spectively, and an inclination angle of 60◦–75◦). Shaposhnikov mal H-rich unevolved main-sequence star. It must be a nuclearly et al. (2012), on the other hand, estimated MBH = 20 ± 3 M evolved object, such as produced by the above mentioned mod- by using an empirical relation between low-frequency QPO and els of Pylyser & Savonije (1988)andPfahletal.(2003). We X-ray spectral shape. The difference in mass estimates could are sure that in MAXI J1659−152 mass transfer from the com- be explained by taking into account the spin of the black hole panion is taking place, because also during quiescence, at least (Kennea et al. 2011; Yamaoka et al. 2012). after the outburst, there is some X-ray emission (see Sect. 1). Because of the rather large uncertainty in possible black-hole This indicates that there is always an accretion disk present in masses, we first assume that the black hole has a mass of MBH the system, which implies that mass transfer continues during 3 M. Using the companion mass estimate from Sect. 4.2.1.1, quiescence. The companion must thus be filling its Roche lobe. this leads to a mass ratio of q = M2/MBH 0.065 (see also The outburst of the system is, therefore, most likely due to some Kennea et al. 2011). Again assuming the companion is Roche- disk instability mechanism (e.g., Lasota 2001), and not due to lobe filling, we can use the relation between the orbital separa- some mass-loss event from a companion that is not filling its tion, a,andq (Eggleton 1983), to get a 1.33 R. The duration Roche lobe. of the ingress and egress of the absorption dips (a few seconds) We estimate the initial mass of the companion star as fol- and the duration of the dip activity (up to 40 min) provide an esti- lows. For the companion star to become evolved in a Hubble mate of the extent of the object being absorbed and the absorber time, its mass must at least have been 1 M. On the other hand, itself, respectively (see, e.g., Kuulkers et al. 1998, and refer- in order for the system to evolve to a short orbital period when ences therein). Assuming the black-hole mass follows the distri- its companion is driving mass transfer (partly) by nuclear evo- bution of known BHXB masses, i.e., taking MBH = 8 M (Özel lution, its companion star cannot have been more massive than et al. 2010; Kreidberg et al. 2012), and following Kuulkers et al. about 1.5 M. The reason for this is that, in order to have the (1998), we find that the upper limit on the extent of the object be- orbital period decrease whilst the companion is evolving, the or- ing absorbed is about 0.002 R or 0.005 R (i.e., about 1000 km bital expansion due to the mass transfer driven by the evolution and 3500 km, respectively), if the absorbing medium corotates of the companion (transfer from the less massive to the more with the binary frame or corotates with matter in the accretion massive star makes the orbit expand) must be more than com- disk, respectively. Similarly, the size of the absorbing medium is pensated by the orbital shrinking due to the orbital angular mo- estimated to be 1.8 R and 6.0 R, respectively. For lower black- mentum loss from the system by magnetic braking. In fact, this hole masses the estimated sizes are somewhat smaller. The size implies that, the system must have started out with an orbital pe- of the object being absorbed is considerably larger than the in- riod below the so-called “bifurcation period”, which is between nermost disk radius (∼30–100 km, Muñoz-Darias et al. 2011; 0.5 days and 0.8 days, depending on the system parameters (see Yamaoka et al. 2012). The X-ray emission thus clearly comes Pylyser & Savonije 1988; Pfahl et al. 2003). Above this limit- from the inner part of the accretion disk, which is possibly sur- ing period in systems with an evolving companion and magnetic rounded by a disk wind or corona. The fact that the absorption braking, the orbital period increases in time, whereas below this dips do not drop to zero intensity at minimum is consistent with critical period it will decrease in time, since the angular momen- this: part of the disk wind or corona stays always visible. The tum loss by magnetic braking wins from the orbital expansion estimated sizes of the absorbing medium are large, i.e., ∼1–3 due to the companion’s internal nuclear evolution. In order to times the orbital separation. This is very unlikely, so we suggest

A32, page 14 of 20 E. Kuulkers et al.: MAXI J1659−152 the absorbing medium to be indeed spread over the outer part of The source reported by Kong et al. (2010) and Kong (2012) the accretion disk, possibly along the accretion disk rim or the seems brighter than expected, which cannot be explained by the region above that, as discussed in Sect. 4.1.1. inclination (see above), a disk contribution or reddening (see Sect. 4.4) alone, and the fact that not all colours are consistent 4.3. Optical counterpart of MAXI J1659−152 with a M5V star, which can also not be solved by the presence of an accretion disk. We, therefore, conclude that we still (Kuulkers Using the fact that for LMXB transients the outburst amplitude et al. 2012a; see also Kennea et al. 2011) can not exclude the Δ = − ( V Vmin Vmax) is related to the orbital period, Porb (Shahbaz possibility that the optical source is a foreground star. & Kuulkers 1998), we derive that the expected outburst amplitude for MAXI J1659−152, is ΔV = 11.4 ± 0.8 mag. The ob- − served optical magnitude during outburst maximum was, Vmax 4.4. Distance to MAXI J1659 152 16.5 mag (see Russell et al. 2010;Kenneaetal.2011). This gives − = ± The distance to MAXI J1659 152 can be estimated in various an expected V-magnitude in quiescence: Vmin 27.9 0.8mag ways. At maxima during outbursts, the optical brightness is dom- (see also Kennea et al. 2011). Applying the correction for in- inated by emission from the accretion disk. Assuming that all of clination on the observation of Vmax (see Sect. 4.4,Eq.(1)) this the optical flux in quiescence comes from the companion, we can would lead to expected V-magnitudes between 27.5 and 26.2, ◦ ◦ use Eq. (5) of Shahbaz & Kuulkers (1998) to derive an estimate for inclination angles between 65 and 80 , respectively. If true, of the absolute disk brightness, which is then only a function it may therefore not be easy to find the optical counterpart in of P . This leads to a rough estimate of M = 1.0±0.8mag. quiescence. orb V,disk To estimate the interstellar reddening, AV , we use the relation A Pan-STARRS 1 (PS1) 3Pi sky-survey observation on from Güver & Özel (2009) between the hydrogen column den- 2010 June 19, however, revealed an optical source consistent sity, NH,andAV . The measured values of NH by the Swift/XRT with the position of MAXI J1659−152 with an AB magnitude during and after the outburst vary between 2.4 and 6 × 1021 cm−2 of about 22.8 in the rP1-band. The source was not detected in the (Kennea et al. 2011; Yamaoka et al. 2012). This is slightly higher other filters (Kong et al. 2010; Kong 2012). Another observation, than the estimated Galactic H I column density in the direc- on 2012 March 23, with the Canada France Hawaii Telescope − tion of MAXI J1659 152 using Kalberla et al. (2005), i.e., the 21 −2 (CFHT), showed the source at a magnitude r 23.7 (Kong weighted average NH = 1.74 × 10 cm , indicating there is in- 2012). trinsic absorption in the system (Kalamkar et al. 2011; Kennea Pan-STARRS uses filters similar to the SDSS g , r , i , et al. 2011). This is not unexpected, given the high inclina- and z filters (plus yP1 and wP1 filter), described in detail by tion of the source (see Sect. 4.1.1). To derive AV we, there- 10 Tonry et al. (2012). The Pan-STARRS system, like SDSS, is fore, use the estimated Galactic NH, which leads to AV 0.8 . = an AB system (e.g., Frei & Gunn 1994), so rP1 r (AB). To With the observed maximum V-band magnitude of Vmax = 16.5 get an estimate of the brightness of the PS1 candidate in the (see Sect. 4.3), using the distance modulus (see Eq. (10) of Johnson V-band, we use the rP1 band detection and the gP1 and Shahbaz & Kuulkers 1998), we then infer that the distance to iP1 upper limits. The limiting magnitudes of the PS1 3Pi sky- MAXI J1659−152 (see also Kennea et al. 2011) is 8.6 kpc, with survey in the gP1 and iP1 bands are estimated to be 23.24 and an estimated uncertainty of 3.7 kpc. Using this distance, we find 22.59 mag, respectively (for a 5-σ point source per visit, e.g., a height above the Galactic plane, z,of2.4± 1.0 kpc (see also Chambers 2006). The Pan-STARRS measurements thus lead to Kennea et al. 2011). gP1−rP1 0.44 and rP1−iP1 0.21. Using the conversions from As noted by Miller-Jones et al. (2011), our derived dis- the Pan-STARRS system to the Johnson system (Tonry et al. tance may be an overestimate, since the relation from Shahbaz 2012), we derive V 23.0 mag. This is consistent with the & Kuulkers (1998) does not take into account the inclination lower limit based on the non-detection in the USNO-B cata- angle of the system. At large inclination angles the projected logue, V > 21 mag, see Kennea et al. (2011), although it is off area of the disk may be smaller, and, therefore, the apparent by about 3 mag from that expected (V 26.2 mag, see above). disk brightness be less. One thus has to correct the observed Another possibility, although we regard it as unlikely, is that the optical magnitude at maximum for this effect. For an optically Shahbaz & Kuulkers (1998) relation breaks down at short orbital thick accretion disk this correction amounts to (Paczynski´ & periods. Schwarzenberg-Czerny 1980, assuming a limb-darkening coef- An M5V star in the Pan-STARRS system gives gP1−rP1 ficient of 0.6, see also Warner 1987): 1.2andrP1−iP1 1.4 (J. Tonry 2012, priv. comm.). Assuming ΔMV = −2.5 log [(1 + 1.5cosi)cosi]. (1) that in quiescence the optical contribution solely comes from the ◦ ◦ companion star, the observed value of rP1−iP1 of the optical can- For inclination angles between 65 and 80 we find that didate is not compatible with a M5V star. Instead, the relation- ΔMV varies from 0.40 to 1.65. This leads to revised dis- ± ± ship between rP1−iP1 and spectral type suggests the candidate to tance estimates between 7.1 3.0 kpc and 4.0 1.7 kpc, respec- be of a type earlier than about G3 (J. Tonry 2012, priv. comm.). tively. Accordingly, z then has values between 2.0 ± 0.9 kpc and Our assumption above may, however, not be correct. The ac- 1.1 ± 0.5 kpc, respectively. cretion disk can in quiescence still contribute significantly (e.g., For a M5V star (Sect. 4.2.1.1) the absolute V-band magni- Jonker et al. 2012, and references therein). Assuming a disk con- tude is about 11.8 (e.g., Zombeck 1990). Although not very contribution of about 50%, the discrepancy between the expected straining, the observed value of V for the proposed optical coun- brightness of the companion and the suggested optical star be- terpart in quiescence (see Sect. 4.3) translates to a lower limit comes about one magnitude less, but is not enough to solve the 10 Note that this is different from the values quoted in D’Avanzo et al. total difference of about three magnitudes. Also, the quiescent (2010; AV 0.34), and Kuulkers et al. (2012a) and Kaur et al. (2012; disk in late-type, short-period BHXBs are not expected to be so A 1.1), based on the measured N from Swift/XRT observations ∼ V H hot ( 6000 K) that the total optical emission mimics a G3-type reported by Kennea et al. (2010). Kennea et al. (2011)usedAV = 1.85 star. The latter can also not be alleviated by the reddening to- (in the UVOT photometric system), which is an upper limit based on a wards MAXI J1659−152, which is only moderate (see Sect. 4.4). value of E(B − V) = 0.606 in the direction of MAXI J1659−152.

A32, page 15 of 20 A&A 552, A32 (2013) to the distance of about 1.1 kpc. Assuming our expected V-band parameters, added a cut-off at 200 keV (which is well outside the magnitude in quiescence (Vmin = 26.2–27.9 mag; Sect. 4.3)we RXTE/PCA energy range), and then integrated between 0.5 keV derive a distance of 5.3–8.7 kpc from the distance modulus. and 10 MeV. To estimate the uncertainty in the bolometric flux Using the PS1 optical counterpart in quiescence proposed by we randomised the spectral parameters using the fit values (and a Kong et al. (2010, Sect. 4.3), and assuming that the companion fixed high-energy cut-off at 200 keV) and their derived 1σ-errors is the sole contributor during optical quiescence, Miller-Jones from Muñoz-Darias et al. (2011b). This was done 100 000 times et al. (2011) estimated the distance to MAXI J1659−152 to and we recorded the resulting integrated fluxes between 0.5 keV be 1.6−4.2 kpc. Assuming the companion to be an M5V (see and 10 MeV. The flux distributions are significantly skewed to- Sect. 4.2.1) or an M2V (Jonker et al. 2012) type star, and us- wards larger fluxes. We, therefore, fitted the flux distributions ing the CFHT optical detection (see Sect. 4.3), Kong (2012) below and above the peak of the distribution with Gaussians derived distances of 2.3–3.8 kpc and 4.6–7.5 kpc, respectively. with different widths. The plus and minus errors in the integrated Kaur et al. (2012) estimate a lower limit of 4 ± 1 kpc, based on flux derived from the spectral fits were then taken as the 1σ the measured radial velocity distribution of the interstellar Na i D widths (see, e.g., Kuulkers et al. 2010a).Wefindanestimated ii +0.42 × −9 −2 −1 and Ca H&K lines. Using the relation between the absolute (unabsorbed) bolometric flux of 4.20−0.34 10 erg cm s . magnitude of a LMXB as a function of the orbital period and the Since the mass of the black hole can be assumed to be at least X-ray luminosity (van Paradijs & McClintock 1994), and assum- 3 M, we derive an upper limit on the transition luminosity for ing that MAXI J1659−152 reached 10% of the Eddington lumi- MAXI J1659−152 of about 10% LEdd. For a canonical black- nosity at maximum of the outburst, Kennea et al. (2011) derived hole mass of 8 M (Özel et al. 2010; Kreidberg et al. 2012) distances in the range 3.2–7.5 kpc for values of AV from 1.85 we derive a transition luminosity of 3.7% LEdd. The latter is to 0. Assuming MBH > 3.2 M and that MAXI J1659−152 at the consistent with the soft to hard state transition luminosity de- peak of the outburst radiates at more than 10% of the Eddington rived by Maccarone (2003), and would suggest a more canoni- limit, they derive a distance >6.1 kpc. Jonker et al. (2012), using cal mass of the black hole in MAXI J1659−152. We note, how- a M2V type companion star and an accretion disk contribution ever, that the black-hole transient GRO J1655−40 did not fit to the quiescent optical light of 50%, derive a distance of 5.9 kpc Maccarone’s (2003) relation. GRO J1655−40 is also a high- with an estimated uncertainty of 2 kpc. Shaposhnikov et al. inclination, dipping source, similar to MAXI J1659−152 (see (2012), using their spectral-timing correlation scaling method, Kuulkers et al. 1998, 2000; Sect. 4.1.1). The above, and given the found an upper limit of 7.6 ± 1.1 kpc, whereas Yamaoka et al. range in distance (Sect. 4.4) and black-hole mass (Sect. 4.2.2) (2012) estimate an upper limit of about 8.6 kpc. The latter au- estimates, leads us to conclude that our test of Maccarone’s thors combined various information, including the fact that the (2003) relation should be used with some caution in the case soft-to-hard transition in BHXB transients occurs at 1–4% of of MAXI J1659−152. the Eddington luminosity (Maccarone 2003; see also Sect. 4.5). The maximum observed 2–10 keV flux during MAXI − − − None of these estimates are particularly robust, however, and J1659–152’s outburst (9 × 10 9 erg cm 2 s 1 on day 13.1; there is a considerable spread. Our initial estimates in the be- Kennea et al. 2011) translates to a maximum 2−10 keV lumi- − ginning of this subsection are consistent with the above quoted nosity of roughly 8 × 1037 erg s 1. This value is not unusual for values; therefore, in the rest of this paper we adhere to our val- transient BHXBs (see, e.g., Dunn et al. 2010). From the X-ray ues of the distance of 8.6 kpc, and the distance above the Galactic spectral fit results by Yamaoka et al. (2012; their model A), plane of z = 2.4 kpc. one can derive that the extrapolated maximum, unabsorbed, 3−200 keV flux occurred on day 13.0, i.e., 10−8 erg cm−2 s−1. For a black-hole with minimum mass of 3 M this gives an up- 4.5. On the soft to hard state transition luminosity per limit on the maximum outburst 3−200 keV luminosity, Lpeak, and the maximum outburst luminosity of about 23% LEdd; for a black hole with 8 M this would lead to Lpeak 8.5% LEdd. The latter value is more or less as expected Black-hole transient sources transit from the soft state to the hard − state when the source luminosity is about 1–4% of the Eddington from the observed relation between the maximum 3 200 keV luminosity and Porb (Wu et al. 2010). luminosity, LEdd, with a mean value of 1.9 ± 0.2% (Maccarone 2003; see also Dunn et al. 2010, for a discussion). Both the spectral and timing behaviour can be used to determine the exact time 4.6. A class of short-period BHXB transients at high Galactic of transition (e.g., Dunn et al. 2010; Belloni 2010; Muñoz-Darias latitudes? et al. 2011a, and references therein). In the hard state, power- law emission dominates and the power-density spectrum shows It is interesting to note (see also Kennea et al. 2011; Yamaoka a strong noise component. MAXI J1659−152 reached the hard et al. 2012) that the two transient BHXBs with the shortest orbital periods, MAXI J1659−152 and Swift J1753.5−0127, are state between day 39.1 and day 40.1 (Muñoz-Darias et al. ◦ ◦ 2011b). Maccarone (2003)usedacut-off power-law spectrum, both found at high Galactic latitudes (bII = 16.5 and 12.2 , Γ= ff respectively), as well as a third short-period transient BHXB: with spectral index, 1.8, and cut-o energy of 200 keV, in- ◦ tegrated between 0.5 keV and 10 MeV to derive the bolomet- XTE J1118+480 (see Zurita et al. 2008, bII = 62.3 ). Given ric correction. To estimate the bolometric flux near the tran- their distances, the corresponding heights above the Galactic sition of the soft-to-hard state for MAXI J1659−152, we used plane are: z = 2.4 kpc (see Sect. 4.4), 1.1 kpc (Zurita et al. the spectrum observed on day 41.0 from Muñoz-Darias et al. 2008) and 1.6 kpc (Jonker & Nelemans 2004), respectively. This +0.02 11 (2011b), which showed Γ=1.80− . We used their spectral can be compared with the observed Galactic z-distribution of 0.03 the BHXBs, which has an rms z-value of 0.625 kpc (Jonker & 11 Nelemans 2004). Recently, another candidate for such a binary Note that Yamaoka et al. (2012) used the estimated bolometric flux − from an observation of the hard state after the state transition, on day has been put forward: Swift J1357.2 0933 (Casares et al. 2011, II ◦ 44.1. Their observation is about 3 days later than the observation we see also Yamaoka et al. 2012; b = 51 ). This system has a take as being near the transition; the flux had declined by about 20% in M4V star as the companion (Rau et al. 2011; Casares et al. 2011) that time. and is subluminous in the radio (Sivakoff et al. 2011), similar to

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MAXI J1659−152 (see Jonker et al. 2012). However, as noted received a kick of 100–200km s−1, the systems which have to in Sect. 4.1.1, the optical outburst amplitude of ∼6 mag (Rau drag along only a companion of low mass (1.5 M) will get et al. 2011) is not consistent with the short period suggested by the highest space velocities. In that case one expects the systems Casares et al. (2011). with the largest space velocities – and therefore the largest mean Swift J1753.5−0127 shares other similarities than only the distances to the galactic plane – to be found among the BHXBs short orbital period with MAXI J1659−152. The shape of the with the shortest orbital periods. (However, since there are also overall outburst light curve as seen by Swift/BAT (Soleri et al. some black holes that do not receive large kicks, one still would 2010), is very similar to that seen for MAXI J1659−152 (see also also expect some of the systems with very short orbital periods Kennea et al. 2011). Swift J1753.5−0127, however, has not been to be located not far from the Galactic plane). seen to turn off again so far, i.e., pre-outburst 15–50 keV flux The alternative possibility, that the systems with the shortest levels have not been reached. Moreover, after it had reached the orbital periods were kicked out of globular clusters, seems very lowest levels about 200 days after the start of the outburst, it has unlikely. Black holes will be born very early in the life of a glob- been seen to vary on years time scales between about 20 mCrab ular cluster (within the first ten million years). They will then by and 125 mCrab (see, e.g., Fig. 1 of Soleri et al. 2010). It has gravitational interactions with the other cluster stars rapidly sink also been active in the radio wavelengths at a lower level com- to the cluster core, where the most massive objects of the cluster pared to when the source was in outburst (Soleri et al. 2010). will concentrate. Calculations (Kulkarni et al. 1993) show that in Low-level soft X-ray and radio activity has been reported for clusters of high central density the rapid dynamical evolution of MAXI J1659−152 (Kennea et al. 2011;Yangetal.2011a,b; the black-hole population in the cluster core leads to the ejection Yang & Wijnands 2011,b; Miller-Jones et al. 2011; Jonker et al. of nearly all the black holes on a short time scale. Because of 2012). In the hard, 15–50 keV band MAXI J1659−152 is not the virtual absence of low-mass stars in these dense cores, it is (yet) detected, from about 65 days after the start of the out- very unlikely that they will have been able to capture a low-mass burst up to now. But it may be detected later again, if it fol- star before they were kicked out. However, for clusters of inter- lows the same trend as that seen for Swift J1753.5−0127. One mediate density, these authors found that some black holes sur- can compare the main outburst light curve of MAXI J1659−152 vive in the cluster, and that some of these surviving black holes also with that of the BHXB transient XTE J1859+226 (Casella could form a LMXB. However, these BHXBs will stay in the et al. 2004; see also Sect. 4.1.2). A high, soft, flux period is fol- cluster, since during a later phase in the evolution of the cluster, lowed by a jump to a low, hard, one. In XTE J1859+226 the such a massive object cannot be kicked out any more by dynam- main outburst lasted also for about a month, with day to day ical interactions with the cluster stars, because these are now of variations similar to that seen for MAXI J1659−152. The orbital much lower mass than the black hole. It thus seems virtually im- period of XTE J1859+226, however, is not as extreme: 6.6 h possible that systems like MAXI J1659−152 were formed in a (Corral-Santana et al. 2011). globular cluster. Yamaoka et al. (2012) argue that MAXI J1659−152 is a runaway micro-quasar, similar to XTE J1118+480, i.e., kicked out of the Galactic plane into the halo. We consider this indeed 4.7. Summary a plausible possibility for explaining the apparently large dis- We have presented here a very detailed analysis of the X-ray and tances of these short-period systems from the Galactic plane − ∼ for the following reasons (we say here “apparently”, because UV light curves of MAXI J1659 152, obtained over 260 days the distances to these systems are still quite uncertain, caus- in 2010–2011 with MAXI, RXTE, Swift and XMM-Newton.Our ing considerable uncertainty in their distances to the Galactic analysis combined soft and hard energies and time resolutions plane). The available evidence for kicks imparted to the black to create a uniform presentation of the source intensity and timing behaviour. We identified two types of variations in the light holes in their formation events suggests that a sizeable fraction ff of black holes may receive rather large kick velocities at birth: curves, absorption and transition dips, characterised by di ering GRO J1655−40 has an observed excess radial velocity relative to spectral properties. The timing studies of these have led us to the its local rest frame of 112 ± 18 km s−1, and simulations of its evo- following conclusions. lution and formation by Willems et al. (2005), including the effects of kicks, indicate that the most likely kick velocity imparted – The absorption dips occur at the orbital period of the system in its formation was 105 km s−1. Similarly, XTE J1118+480 has and are due to the combined effects of high orbital inclina- an excess radial velocity of 145 ± 25 km s−1 and simulations of tion and obscuration by material from the companion star in- its evolution and formation by Fragos et al. (2007) indicate that teracting with the accretion disk. The presence of these dips a most likely kick velocity of about 200 km s−1 imparted in its has allowed us to measure precisely the orbital period of the formation event. On the other hand, simulations of the evolution binary at 2.414 ± 0.005 h. This is the shortest BHXB period of the Cyg X-1 system by Wong et al. (2012) indicate that the measured to date and is confirmed (Kuroda et al. 2010) with kick imparted at its formation was less than 77 km s−1 and most modulations of the optical light curve of the system. probably not more than about 40 km s−1 (see also Reid et al. – Using the absorption dips we have constrained the inclina- 2011). The BHXB V404 Cyg has a peculiar velocity of about tion of MAXI J1659−152 to be between 65◦–80◦ and the 40 km s−1 (Miller-Jones et al. 2009), indicating that its black hole spectral type of the companion star to be M5V. did not receive a velocity kick larger than this value. If many – We also identified transition dips, which are most likely a black holes would receive a velocity kick of the same order as result of instabilities in the inner accretion flow. During these those of GRO J1655−40 and XTE J1118+480, i.e., about 100 dips the source softens with increasing intensity, in contrast to 200 km s−1, one may expect that many of the systems with to the absorption dips, which are less energy dependent with short orbital periods will have high runaway velocities. As men- increasing count rates. The exact origin of these dips remains tioned in Sect. 4.2.1, in order for a BHXB to evolve to a very unknown. short orbital period, the initial mass of the companion star must – Using an estimate of the black hole mass of >3 M,anda have been 1.5 M. If a typical black hole of mass 6 to 10 M ratio between the compact object and the companion mass

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of q = M2/MBH 0.065, we estimate the binary orbital Casella, P., Belloni, T., Homan, J., & Stella, L. 2004, A&A, 426, 587 separation to be a 1.33 R. Chambers, K. C. 2006, Mission Concept Statement for PS1; Pan-STARRS – The very short orbital period of the system allowed us to Telescope #1, University of Hawai‘i at Manoa, Doc. No. PSDC-230-002-00 Corral-Santana, J. M., Casares, J., Shahbaz, T., et al. 2011, MNRAS, 413, L15 successfully argue that the companion is a nuclearly evolved Corral-Santana, J. M., Casares1, J., Muñoz-Darias, T., et al. 2013, Science, 339, star with initial mass of about 1.5 M, which evolved to 0.19 1048 (0.17) M during 4.6 (5.7) billion years (depending on the Cowley, A. P., Crampton, D., Hutchings, J. B., Remillard, R., & Penfold, J. E. evolutionary sequence). 1983, ApJ, 272, 118 ± D’Avanzo, P., Goldoni, P., Patruno, A., et al. 2010, ATel, #2900 – We adopt an inferred distance to the source of 8.6 3.7 kpc de Jong, J. A., van Paradijs, J., & Augusteijn, T. 1996, A&A, 314, 484 and a distance above the Galactic plane of z = 2.4 ± 1.0 kpc, de Ugarte Postigo, A., Flores, H., Wiersema, K., et al. 2010, GCN Circ., 11307 which translates to maximum 2−10 keV and 3−200 keV den Herder, J. W., Brinkman, A. C., Kahn, S. M., et al. 2001, A&A, 365, L7 luminosities of ∼8 × 1037 erg s−1 and ∼9 × 1037 erg s−1, Díaz Trigo, M., Parmar, A. N., Boirin, L., Méndez, M., & Kaastra, J. S. 2006, respectively. A&A, 445, 179 Díaz Trigo, M., Parmar, A. N., Boirin, L., et al. 2009, A&A, 493, 145 – There are by now three BHXB sources with relatively short Done, C., & Diaz Trigo, M. 2010, MNRAS, 407, 2287 periods and high Galactic scale heights. These have been ar- Dunn, R. J. H., Fender, R. P., Körding, E. G., Belloni, T., & Cabanac, C. 2010, gued to be runaway micro-quasars, i.e., systems kicked out MNRAS, 403, 61 of the plane into the halo. We argue that this hypothesis Durant, M., Gandhi, P., Shahbaz, T., Peralta, H. H., & Dhillon, V. S. 2009, − MNRAS, 392, 309 is better supported by the properties of MAXI J1659 152, Eggleton, P. P. 1983, ApJ, 268, 368 when contrasted with the suggestion that the system was Elvis, M., Page, C. G., Pounds, K. A., Ricketts, M. J., & Turner, M. J. L. 1975, formed in a globular cluster. Nature, 257, 656 Evans, P. A., Beardmore, A. P., Page, K. L., et al. 2009, MNRAS, 397, 1177 Future X-ray spectral analysis, as well as additional multi- Foulkes, S. B., Haswell, C. A., & Murray, J. R. 2010, MNRAS, 401, 1275 wavelength studies, both in outburst and quiescence, will elu- Fragos, T., Willems, B., Ivanova, N., & Kalogera, V. 2007, AIPC, 924, 673 Frank, J., King, A. R., & Lasota, J.-P., 1987, A&A, 178, 137 cidate the shortest period BHXB further. Frei, Z., & Gunn, J. E. 1994, AJ, 108, 1476 Gehrels, N., Chincarini, G., Giommi, P., et al. 2004, ApJ, 611, 1005 Acknowledgements. Partly based on observations obtained with XMM-Newton, Greiner, J., Rau, A., & Schady, P. 2012, ATel, #4030 an ESA science mission with instruments and contributions directly funded by Güver, T., & Özel, F. 2009, MNRAS, 400, 2050 ESA Member States and NASA. This research has made use of data obtained Hill, J. E., Burrows, D. N., Nousek, J. A., et al. 2004, SPIE, 5165, 217 through the High Energy Astrophysics Science Archive Research Center Online / / Haswell, C. A., King, A. R., Murray, J. R., & Charles, P. A. 2001, MNRAS, 321, Service, provided by the NASA Goddard Space Flight Center. The MAXI GSC 475 data are provided by RIKEN, JAXA and the MAXI team, whilst the Swift/BAT / Homan, J., & Belloni, T. 2005, Ap&SS, 300, 107 transient monitor results are provided by the Swift BAT team. We especially Homan, J., Wijnands, R., van der Klis, M., et al. 2001, ApJS, 132, 377 thank the XMM-Newton Science Operations Centre for their prompt schedul- Horne, K. 1985, MNRAS, 213, 129 ing of the Target of Opportunity observations, 5 h between trigger and start of Howell, S. B., Nelson, L. A., & Rappaport, S. A. 2001, ApJ, 550, 897 observation on September 27! We would also like to thank the Swift and RXTE Jahoda, K., Markwardt, C. B., Radeva, Y., et al. 2006, ApJS, 163, 401 teams for their scheduling of the many monitoring observations. The research Jansen, F., Lumb, D., Altieri, B., et al. 2001, A&A, 365, L1 leading to these results has received funding from the European Community’s / / Johnson, H. L. 1966, ARA&A, 4, 193 Seventh Framework Programme (FP7 2007 2013) under grant agreement num- Jonker, P. G., & Nelemans, G. 2004, MNRAS, 354, 355 ber ITN 215212 “Black Hole Universe”. TMB acknowledges support to ASI- Jonker, P., Miller-Jones, J. C. A., Homan, J., et al. 2012, MNRAS, 423, 3308 / / / INAF grant I 009 10 0, as well as funding via an EU Marie Curie Intra-European Kahn, D. A. 2010, GCN Circ., 11299 Fellowship under contract no. 2011-301355. EK thanks John Tonry for discus- Kalamkar, M., Homan, J., Altamirano, D., et al. 2011, ApJ, 732, L2 sions regarding the Pan-STARRS 1 optical candidate, Vik Dhillon for supplying Kalberla, P. M. W., Burton, W. B., Hartmann, D., et al. 2005, A&A, 440, 775 the “PERIOD” analysis package, which we used in our periodicity analysis, and Kaur, R., Kaper, L., Ellerbroek, L. E., et al. 2012, ApJ, 746, L23 Kazutaka Yamaoka for providing the estimated 3–200 keV fluxes from RXTE Kennea, J. A., Krimm, H., Mangano, V., et al. 2010, ATel, #2877 spectral fits. E.K. and T.M.D. acknowledge Sara Motta for her comments on the Kennea, J. A., Romano, P., Mangano, V., et al. 2011, ApJ, 736, 22 RXTE spectral analysis. E.K. and M.D.T. thank Roberto Vio for a discussion on Kennea, J. A., Altamirano, D., Evans, P. A., et al. 2012a, ATel, #4034 the periodograms. J.C. was supported by ESA-PRODEX contract N: 90057. Last / Kennea, J. A., Yang, Y. J., Altamirano, D., et al. 2012b, ATel, #4044 but not least, we thank the referee for his her careful reading of the manuscript. Kepler, J. 1619, in Harmonices Mundi, Libri V, Lincii Austriae Knigge, C., Baraffe, I., & Patterson, J. 2011, ApJS, 194, 28 Note added in proof. In a recent paper, Corral-Santana et al. (2013)de- Kong, A. K. H. 2012, ApJ, 760, L27 termined the orbital period of Swift J1357.2-093313 to be 2.8 ± 0.3h. Kong, A. K. H., Lin, C.-C, Chen, Y. T., et al. 2010, ATel, #2976 With a distance of ∼1.6 kpc it is located at ∼1.2 kpc above the Galactic Kong, A. K. H., Yang, Y. J., & Wijnands, R. 2011, ATel, #3524 plane. This places the source indeed among the BHXB sources with Kreidberg, L., Bailyn, C. D., Farr, W. M., & Kalogera, V. 2012, ApJ, 757, 36 relatively short orbital periods and high Galactic scale heights. 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A32, page 19 of 20 A&A 552, A32 (2013)

Table 1. Log of X-ray observations with RXTE, Swift,andXMM-Newton during the main outburst of MAXI J1659−152 presented in this paper, ordered along the start time of the observation.

Day1 Start time (UT) Exp.2 ObsID Satellite Day1 Start time (UT) Exp.2 ObsID Satellite 0.3 2010-09-25 07:49 19 667 00434928000 Swift 16.6 2010-10-11 15:30 2028 95108-01-25-00 RXTE 1.0 2010-09-26 00:07 16 214 00434928001 Swift 17.0 2010-10-12 00:36 3348 95108-01-26-00 RXTE 1.6 2010-09-26 13:31 9909 00434928002 Swift 17.7 2010-10-12 16:27 2453 95108-01-27-00 RXTE 2.0 2010-09-27 00:13 19 108 00434928003 Swift 18.4 2010-10-13 09:43 2697 95108-01-28-00 RXTE 2.7 2010-09-27 16:15 51 916 0656780601 XMM-Newton 18.8 2010-10-13 19:24 1284 00434928023 Swift 3.0 2010-09-28 00:53 16 910 95358-01-02-00 RXTE 19.5 2010-10-14 11:37 1705 00434928025 Swift 3.3 2010-09-28 07:06 10 009 00434928005 Swift 19.9 2010-10-14 21:32 3516 95108-01-30-00 RXTE 4.1 2010-09-29 01:58 1706 95358-01-02-01 RXTE 20.0 2010-10-15 00:28 1240 00434928026 Swift 4.2 2010-09-29 05:14 2318 00434928007 Swift 20.2 2010-10-15 05:27 3444 95118-01-01-00 RXTE 5.1 2010-09-30 02:00 1410 95358-01-02-02 RXTE 20.6 2010-10-15 13:27 1390 00434928027 Swift 5.2 2010-09-30 05:37 2719 00434928008 Swift 20.7 2010-10-15 16:24 3554 95118-01-01-01 RXTE 6.2 2010-10-01 05:45 3369 95358-01-03-00 RXTE 21.0 2010-10-16 00:34 1305 00434928028 Swift 6.2 2010-10-01 05:45 2594 00434928009 Swift 21.1 2010-10-16 03:25 3374 95118-01-02-00 RXTE 6.4 2010-10-01 10:44 2130 95108-01-02-00 RXTE 21.5 2010-10-16 11:45 1325 00434928029 Swift 7.1 2010-10-02 02:31 1763 95358-01-03-01 RXTE 22.0 2010-10-17 00:38 1390 00434928030 Swift 7.1 2010-10-02 02:37 2369 00434928010 Swift 22.2 2010-10-17 04:25 829 95118-01-03-01 RXTE 7.5 2010-10-02 11:50 1905 95108-01-03-00 RXTE 22.8 2010-10-17 18:55 2464 95118-01-03-00 RXTE 7.8 2010-10-02 18:22 1175 95108-01-04-00 RXTE 23.0 2010-10-18 00:46 1289 00434928032 Swift 8.1 2010-10-03 01:41 994 95108-01-05-00 RXTE 23.7 2010-10-18 16:36 3556 95118-01-04-00 RXTE 8.1 2010-10-03 02:36 1720 00434928011 Swift 24.0 2010-10-19 00:24 3027 95118-01-05-00 RXTE 8.2 2010-10-03 04:50 3352 95358-01-03-02 RXTE 24.0 2010-10-19 00:50 1290 00031843001 Swift 8.5 2010-10-03 11:20 1737 95108-01-06-00 RXTE 24.6 2010-10-19 13:31 1245 00031843002 Swift 8.9 2010-10-03 21:03 1378 95108-01-07-00 RXTE 24.9 2010-10-19 20:51 2358 95118-01-05-01 RXTE 9.1 2010-10-04 02:40 3314 00434928012 Swift 25.0 2010-10-20 00:56 1290 00031843003 Swift 9.1 2010-10-04 02:43 3352 95108-01-08-00 RXTE 25.3 2010-10-20 06:17 3495 95118-01-06-00 RXTE 9.5 2010-10-04 10:52 2216 95108-01-09-00 RXTE 25.5 2010-10-20 12:21 1089 00031843004 Swift 9.7 2010-10-04 17:19 1309 95108-01-10-00 RXTE 25.7 2010-10-20 17:16 3468 95118-01-06-01 RXTE 10.1 2010-10-05 02:47 3414 00434928013 Swift 26.0 2010-10-21 00:53 985 00031843005 Swift 10.6 2010-10-05 13:43 1847 95108-01-11-00 RXTE 26.1 2010-10-21 02:40 2883 95118-01-07-01 RXTE 10.8 2010-10-05 20:02 1676 95108-01-12-00 RXTE 26.6 2010-10-21 13:46 985 00031843006 Swift 11.1 2010-10-06 02:53 3284 00434928014 Swift 26.7 2010-10-21 16:48 3214 95118-01-07-00 RXTE 11.4 2010-10-06 09:48 2397 95108-01-13-00 RXTE 27.0 2010-10-22 00:37 893 95118-01-08-00 RXTE 11.7 2010-10-06 17:55 1443 95108-01-14-00 RXTE 27.0 2010-10-22 00:57 1264 00031843007 Swift 12.1 2010-10-07 01:17 3409 95108-01-15-00 RXTE 27.2 2010-10-22 05:51 1040 00031843008 Swift 12.1 2010-10-07 02:47 1575 00434928015 Swift 27.8 2010-10-22 19:22 1940 95118-01-09-00 RXTE 12.4 2010-10-07 09:28 2337 95108-01-16-00 RXTE 29.2 2010-10-24 05:54 1162 95118-01-10-00 RXTE 12.7 2010-10-07 15:38 1595 00434928016 Swift 30.2 2010-10-25 05:24 888 95118-01-11-00 RXTE 12.7 2010-10-07 15:47 1602 95108-01-17-00 RXTE 31.0 2010-10-26 00:39 868 95118-01-12-00 RXTE 13.0 2010-10-07 23:37 858 95108-01-18-00 RXTE 32.5 2010-10-27 12:25 2197 95118-01-13-00 RXTE 13.0 2010-10-08 00:00 1015 95108-01-18-01 RXTE 33.5 2010-10-28 11:54 774 95118-01-14-00 RXTE 13.1 2010-10-08 02:52 1630 00434928017 Swift 34.5 2010-10-29 11:31 1152 95118-01-15-00 RXTE 13.7 2010-10-08 16:55 1662 95108-01-19-00 RXTE 35.3 2010-10-30 07:54 1412 95118-01-15-01 RXTE 14.0 2010-10-09 00:49 1900 95108-01-20-00 RXTE 36.3 2010-10-31 07:24 1458 95118-01-16-00 RXTE 14.1 2010-10-09 02:58 1584 00434928019 Swift 37.2 2010-11-01 05:20 1467 95118-01-16-01 RXTE 14.5 2010-10-09 11:39 2309 95108-01-21-00 RXTE 38.0 2010-11-02 00:14 1715 95118-01-17-00 RXTE 14.7 2010-10-09 15:49 1595 00434928020 Swift 39.0 2010-11-03 01:11 1457 95118-01-17-01 RXTE 15.1 2010-10-10 01:27 1009 00434928021 Swift 40.1 2010-11-04 01:15 786 95118-01-18-00 RXTE 15.1 2010-10-10 03:05 3507 95108-01-22-00 RXTE 41.0 2010-11-05 00:16 2348 95118-01-19-00 RXTE 15.5 2010-10-10 12:42 1629 00434928022 Swift 42.2 2010-11-06 04:32 1139 95118-01-20-00 RXTE 15.7 2010-10-10 15:54 1914 95108-01-23-00 RXTE 44.1 2010-11-08 01:58 1608 95118-01-21-00 RXTE 16.2 2010-10-11 04:03 3513 95108-01-24-00 RXTE

Notes. (1) Day = MJD – 55 464, where MJD 55 464 corresponds to UT 2010, September 25, 0:00. (2) Total on-source exposure time in s.

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